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<div></div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/White-light_Emission_and_Non-thermal_ElectronsWhite-light Emission and Non-thermal Electrons2018-12-07T20:14:37Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = White-light Emission and Non-thermal Electrons <br />
|number = 334<br />
|first_author = Kyoung-Sun LEE <br />
|second_author = <br />
|publish_date = 8 October 2018<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::335]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::333]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
A [https://www.britannica.com/science/solar-flare solar flare]<br />
produces emissions detected all across the electromagnetic spectrum,<br />
and thus involving all "layers" of the solar atmosphere.<br />
In particular the <br />
[https://solarscience.msfc.nasa.gov/chromos.shtml chromosphere] <br />
may contain some of the decisive physics, based upon the newest data.<br />
<br />
Recent space-borne UV - EUV - X-ray spectroscopic observations allow<br />
us to investigate the flare plasma properties and their temporal<br />
evolution throughout the solar atmosphere, chromosphere to the<br />
corona. <br />
For example, the <br />
[https://en.wikipedia.org/wiki/Christian_Doppler Doppler]<br />
[https://www.physicsclassroom.com/class/waves/Lesson-3/The-Doppler-Effect velocity]<br />
from the UV-EUV spectra gives information about plasma flows in the <br />
solar atmosphere which can be compared with the velocities from <br />
simulation models, helping to understand the flare heating process [Ref. 1]. <br />
Our diagnostic abilities for temperature and density have been <br />
greatly extended recently with wide wavelength coverage of the <br />
space-borne spectroscopy and simulations [Ref. 2]. <br />
Taking advantage of the combining UV-EUV imaging spectroscopy,<br />
in this Nugget, we report the flare plasma dynamics and energy flux<br />
through the solar atmosphere during a <br />
[https://en.wikipedia.org/wiki/Solar_flare white-light flare] (WLF) using the<br />
multiple spectroscopic observations, and then, we try to understand<br />
how the visible continuum is produced [Ref. 3].<br />
<br />
== A flare kernel observed by SDO, IRIS, Hinode, and RHESSI ==<br />
<br />
[https://www.solarmonitor.org/?date=20141022 NOAA AR 12192] <br />
was the largest active region in <br />
[https://en.wikipedia.org/wiki/Solar_cycle_24 Solar Cycle 24].<br />
Many coordinated observations concentrated on this region, and an<br />
[http://www.thesuntoday.org/missions/goes/ X-class flare] was observed <br />
by multiple space-based observations covering the whole flare duration <br />
(SOL2014-10-22T14:02). <br />
Figure 1 shows the flare<br />
observations from <br />
[https://www.swpc.noaa.gov/products/goes-x-ray-flux GOES], <br />
[http://hmi.stanford.edu SDO/HMI] continuum, <br />
[https://en.wikipedia.org/wiki/Solar_Dynamics_ObservatorySDO/AIA], <br />
[http://iris.lmsal.com IRIS], <br />
[http://global.jaxa.jp/projects/sat/solar_b/ Hinode],<br />
and RHESSI spacecraft. <br />
These multi-wavelength spectroscopic observations show<br />
that the flare kernel was localized during the impulsive phase and<br />
that the HXR emission, chromospheric intensity, and WL continuum emission<br />
in the kernel are spatially and temporally correlated. <br />
Figure 2 shows the Doppler velocity and line width of UV-EUV spectra in the<br />
localized bright kernel with a time and a peak formation temperature.<br />
As Figure 2 shows, we found that strong evaporation flows (up-flows<br />
in hot lines) occurs and there are strong enhancements<br />
in line width during the first peak of the HXR emission, which is<br />
also coincident with the timing of the WLF. <br />
This may indicate that<br />
electron beam heating produces strong evaporation flows and there<br />
is turbulence from the reconnection or non-thermal electron heating.<br />
<br />
[[File:334f1.png|600px|thumb|center|Figure 1:<br />
Context images of the X1.6 flare on 2014 October 22 14:06 UT.<br />
The arrows indicate the bright kernel we studied. <br />
<br><br />
(a) The GOES X-ray light curve (top), its time derivative (middle), and the <br />
RHESSI count rates for the different energy bands (bottom). <br />
The vertical dashed lines mark the times of the three peaks in the<br />
time derivative of the GOES X-ray light curve. <br />
<br><br />
(b) SDO/AIA 1600&Acirc; image at the timing of the first hard<br />
X-ray peak (~14:06 UT). <br />
<br><br />
(c) SDO/HMI continuum and a running difference image at 14:05:45 UT. <br />
<br><br />
(d) IRIS C II 1330 slit jaw images overlaid<br />
with EIS 195 intensity contours (gray line) and the location of<br />
IRIS raster (white box) at 14:06 UT. <br />
<br><br />
(e) HMI continuum difference<br />
image with HXR (30100 keV) and SXR (1225 keV) contours overlaid<br />
from the RHESSI image covering 14:05:3214:06:32 UT. <br />
Red and<br />
green contours correspond to 50%, 60%, 70%, 80%, and 90% of the HXR<br />
and SXR intensity, respectively.<br />
]]<br />
<br />
[[File:334f2.png|600px|thumb|center| Figure 2: <br />
(a) Temporal variation of the Doppler velocity (top) and line width<br />
(bottom) for the bright kernel from the EIS spectral lines: He II,<br />
Fe XII, Fe XV, and Fe XXIII. The vertical dashed lines correspond<br />
to the observing times of the panel (b). <br />
<br><br />
(b) Doppler velocities<br />
from EIS and IRIS as a function of the peak formation temperature<br />
for the bright kernel at different times, the impulsive phase (top)<br />
and the gradual phase (bottom). Diamonds indicate the Doppler<br />
velocities from the IRIS spectra while crosses represent the Doppler<br />
velocities from the EIS spectra. Black and red indicate the velocities<br />
calculated from the single and multiple Gaussian components relative<br />
to the rest wavelengths.<br />
]]<br />
<br />
== White light emission produced by non-thermal electrons ==<br />
<br />
The continuum enhancement of the flare kernel (WLF) implies that<br />
the flare energy is transported to the lower atmosphere and produces<br />
the chromospheric and photospheric emission. <br />
To understand the heating process and how the WLF kernel is produced, <br />
we tried to estimate the deposited and dissipated energies from the X-ray, UV,<br />
and continuum intensities. <br />
Figure 3 shows the IRIS <br />
[https://ned.ipac.caltech.edu/level5/Ewald/Grotrian/frames.html Mg II] <br />
spectra in the flare kernel and the temporal variation of the<br />
estimated energy flux from the X-ray, UV, and continuum intensity.<br />
The Mg II subordinate line (Mg II triplet) is in emission during<br />
the impulsive phase of the flare while the Mg II triplet mostly was<br />
observed as absorption in quiet regions (Figure 3a). <br />
It has been reported that a strong enhancement of the intensity <br />
ratio of the Mg II core to wing (emission) implies the existence of a steep<br />
temperature gradient and heating [Refs. 2, 4]. <br />
Taking the measured energy flux from this spectra as an estimate of the <br />
amount of energy dissipated in the chromosphere, it is about 6-22% of <br />
the energy deposited by the accelerated non-thermal electrons as measured by<br />
RHESSI, assuming the cutoff energy of 30-40 keV (Figure 3b). <br />
This result implies that the majority of the energy from the non-thermal<br />
electrons accelerated in the corona is still available to directly<br />
produce a WLF in this event. <br />
The correlated temporal variations of the HXR, WL, explosive evaporation <br />
flows, and the Mg II line response,<br />
together with the comparison of the energy flux through the corona<br />
to the chromosphere, imply that the flare heating and evaporation<br />
flow are driven by non-thermal electrons, though we cannot rule out<br />
a possible contribution from <br />
[https://en.wikipedia.org/wiki/Alfvén_wave Alfv&eacute;n wave] heating.<br />
<br />
[[File:334f3.png|600px|thumb|center| Figure 3: <br />
(a) IRIS detector images of the Mg II h & k spectral windows overlaid<br />
with the averaged spectral line profiles at the location marked by<br />
the two horizontal green lines. The solid lines represent the line<br />
profiles during the pre-flare (top) and impulsive phase (bottom)<br />
and dotted lines correspond the line profile before the flare (around<br />
12:00 UT). The arrows indicate the emission of the Mg II triplet<br />
lines. <br />
<br><br />
(b) The energy flux of the bright kernel during the flare<br />
impulsive phase estimated from the RHESSI HXR emission with different<br />
threshold energies of 30 keV (solid line), 40 keV (dashed line),<br />
and 50 keV (dotted line), the Mg II triplet intensity observed by<br />
IRIS, and the WL continuum emission from SDO/HMI.<br />
]]<br />
<br />
== Conclusions ==<br />
<br />
We have investigated the temporal variation of the spectral properties<br />
of a WLF kernel and estimated the energy flux at different wavelengths.<br />
By comparison of the measured spectral properties and energy<br />
fluxes in different wavelengths, it appears that the flare emissions<br />
could have been directly produced by non-thermal electrons. <br />
The flare we investigated is quite strong (X1.6 classs), and while <br />
the accelerated non-thermal electrons from this strong flare could <br />
produce a WLF directly in this event, we still do not know whether <br />
most WLFs can be produced by the energy from the accelerated electrons. <br />
We particularly do not know yet whether this process<br />
works even in small flares. <br />
Therefore, further studies applying<br />
similar techniques to other flares that produce visible continuum emission are<br />
required to confirm whether they can also be produced directly or<br />
not.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2016ApJ...816...89P "Simultaneous IRIS and Hinode/EIS Observations and Modeling of the 2014 October 27 X2.0 Flare"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2015ApJ...806...14P "The Formation of IRIS Diagnostics. IV. The Mg II Triplet Lines as a New Diagnostic for Lower Chromospheric Heating"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2014Sci...346C.315P "Hot explosions in the cool atmosphere of the Sun"]<br />
<br />
[4][http://adsabs.harvard.edu/abs/2017ApJ...836..150L "IRIS, Hinode, SDO, and RHESSI Observations of a White Light Flare Produced Directly by Nonthermal Electrons"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/CORONAS/SPIRIT_Mg_XII_and_NanoflaresCORONAS/SPIRIT Mg XII and Nanoflares2018-12-07T20:14:24Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = CORONAS/SPIRIT Mg XII and nanoflares <br />
|number = 335<br />
|first_author = Anton Reva <br />
|second_author = <br />
|publish_date = 22 October 2018<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::336]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::334]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
The high temperature of the solar corona has several possible explanations,<br />
and correspondingly intense theoretical and observational work trying <br />
to identify the right one(s).<br />
In nanoflare heating theory, the corona is heated by a large number of<br />
small-scale flare-like events (nanoflares). <br />
If the nanoflare frequency is low, then some hot plasma should always exist <br />
in non-flaring active regions, since the apparently steady emission<br />
would reflect the lower average temperature of the nanoflare events.<br />
The detection of hot plasma in non-flaring active regions would therefore be<br />
indirect evidence of low-frequency nanoflare heating (a <br />
[https://en.wikipedia.org/wiki/Chekhov%27s_gun "smoking gun"]). <br />
The absence of the hot emission can help to constrain the nanoflare frequency.<br />
The breadth of the distribution of emission measure if it is not<br />
isothermal, i.e. the "differential emission measure" (DEM), thus reflects<br />
the presence or absence of nanoflares.<br />
<br />
It is hard to find faint hot-plasma emission with conventional<br />
non-monochromatic imagers (like [https://aia.lmsal.com AIA] or [https://xrt.cfa.harvard.edu XRT]), because their images<br />
contain broad temperature contributions. <br />
This Nugget describes an upper limit on the hot-plasma DEM using <br />
direct observations of the hot plasma by the Mg XII spectroheliograph <br />
SPIRIT (Ref. [1]) onboard the<br />
[https://space.skyrocket.de/doc_sdat/koronas-foton.htm CORONAS-F]<br />
satellite observatory.<br />
This instrument imaged coronal hot plasma without any low-temperature <br />
background contamination. <br />
We compare the obtained limit<br />
with the result of recent numerical simulations and aim to place<br />
constraints on the parameters of the nanoflare-heating model <br />
(Refs. [2,3]).<br />
<br />
== Experimental Data ==<br />
<br />
The Mg XII spectroheliograph obtained monochromatic images of the<br />
solar corona in the Mg XII 8.42 &Acirc; line. <br />
This line emits mainly at temperatures higher than 4 MK. <br />
The Mg XII images do not contain a solar limb (hence no contribution<br />
from the general corona) or any other low-temperature background.<br />
<br />
We studied the period from 18-28 February 2002. At this time, the<br />
Mg XII spectroheliograph worked with a 105-second cadence almost<br />
without data gaps. In these observations, the Mg XII spectroheliograph<br />
registered binned images with a spatial resolution of 8 arc seconds<br />
and a 37-second exposure time.<br />
<br />
== Results ==<br />
<br />
The Mg XII data (Figure 1) shows that only two types of hot objects<br />
were present on the Sun: the first are small isolated flare-like<br />
phenomena (Ref. [4]). During the period of observations, they occurred at<br />
a rate of about 20 per day. <br />
The second are large hot structures inside active regions. <br />
These are produced during flares or sequences of<br />
microflares. <br />
Their X-ray emission is highly variable. After the flare ended, these large structures faded away.<br />
<br />
[[File:335f1.png|600px|thumb|center|Figure 1:<br />
Hot plasma observed by the Mg XII spectroheliograph.<br />
Left: EIT image. Right: Mg XII spectroheliograph image (blue and<br />
green correspond to low intensities, red and yellow to high<br />
intensities). 1) Small flare-like objects. 2) Flaring active region<br />
with hot plasma. 3) Non-flaring active region without hot plasma.<br />
]]<br />
<br />
Except for rare microflares, there was no hot plasma in non-flaring<br />
active regions. Below we discuss this fact and try to estimate the<br />
upper limit of the hot-plasma emission measure and nanoflare<br />
frequency, which may be consistent with this lack of signal.<br />
<br />
To estimate the relative amount of the hot and warm (main temperature<br />
component) plasma, we reconstructed the DEM of the non-flaring<br />
active region [https://www.solarmonitor.org/index.php?date=20020220&region=09833&indexnum=1 NOAA 09833] <br />
(see Figure 2). <br />
We extracted the fluxes<br />
of the selected active region from the Mg XII, EIT 171 &Acirc;, 195 &Acirc;, and<br />
284 &Acirc; channels. Using these four fluxes, we reconstructed the DEM<br />
with a genetic algorithm.<br />
<br />
[[File:335f2.png|600px|thumb|center|Figure 2:<br />
Non-flaring active region NOAA 09833, which was used to determine<br />
the DEM. a) EIT 171 &Acirc; image; b) EIT 195 &Acirc; image; <br />
c) EIT 284 &Acirc; image.<br />
d) Spectroheliograph Mg XII image. The rectangle marks the active<br />
region that was used for DEM measurements<br />
]]<br />
<br />
The amount of plasma with T =5 MK (logT = 6.7) is four orders of<br />
magnitude lower than the main temperature component, and the amount<br />
of plasma with T = 10 MK is four to five orders lower (see Figure<br />
3).<br />
<br />
[[File:335f3.png|600px|thumb|center|Figure 3: The<br />
DEM of the non-flaring active region marked in Figure 2.<br />
Red shows the DEMs obtained during different runs of a genetic<br />
algorithm, and black represents their median. The blue-dashdotted<br />
DEM-loci show EIT 171 &Acirc;, green-dotted DEM-loci are EIT 195 &Acirc;, <br />
red-dashed DEM-loci show EIT 284 &Acirc;, and purple-solid DEM-loci represent <br />
the Mg XII spectroheliograph.<br />
]]<br />
<br />
At log T > 7.0, only the Mg XII flux constrains the DEM. The DEM<br />
in Figure 3 at log T > 7.0 shows the values that could be added to<br />
the DEM without increasing the Mg XII flux to a level greater than<br />
the noise. These values are the DEM upper limit.<br />
<br />
== Comparison with Simulations ==<br />
<br />
Numerical simulations (Ref. [5]) show how a coronal<br />
loop should react to the sequence of nanoflares depending on the<br />
nanoflare frequency. <br />
The author found that for a low nanoflare frequency, the DEM has a <br />
hot component. <br />
The hot component vanishes for a high nanoflare frequency.<br />
For each different regimes and for different delays between nanoflares,<br />
the simulations provide DEMs. <br />
For each of these plots, we manually measured the ratio of the DEM at<br />
10 MK (DEM10) to the DEM of the main temperature component (DEM<sub>max</sub>).<br />
Then we used the obtained values to build a plot of how this ratio<br />
depends on the delay between the nanoflares (see Figure 4).<br />
<br />
Figure 4 shows that the relative amount of hot plasma rapidly<br />
diminishes with the decrease of the delay between nanoflares. The<br />
Mg XII data show that this ratio should be lower than 0.0001. Therefore,<br />
the delay between the nanoflares should be shorter than 500 seconds.<br />
<br />
[[File:335f4.png|600px|thumb|center|Figure 4: The<br />
ratio of DEM10 to the DEM<sub>max</sub> as a function of the average<br />
time between nanoflares. <br />
The data were taken from the simulations of Ref. [5].<br />
The black line denotes the simulations where nanoflares have power-law <br />
energy distribution with a slope m = -2.5 and the delay between <br />
nanoflares is proportional to their energy. <br />
The red line denotes the simulations where nanoflares have<br />
a power-law energy distribution with a slope m = -1.5 and the delay<br />
between nanoflares is proportional to their energy. The blue line<br />
denotes the simulations where nanoflares have a power-law energy<br />
distribution with a slope m = -2.5 and the delay between nanoflares<br />
is fixed. The purple line denotes the Mg XII upper limit on the<br />
ratio of the DEMs of the hot and warm components.<br />
]]<br />
<br />
== Conclusions ==<br />
<br />
The CORONAS-F/SPIRIT data constrain the high-temperature DEM levels in the solar corona, and do not find any evidence for the <br />
"smoking gun" signature expected for nanoflare heating. <br />
Because the instrument is uniquely monochromatic, and because of the extent of the data, these constraints are the most direct ones yet published.<br />
The details of this analysis appear in Ref. [6].<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2003MNRAS.338...67Z "Dynamic 10 MK plasma structures observed in monochromatic full-Sun images by the SPIRIT spectroheliograph on the CORONAS-F mission"]]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2002AstL...28..401O "Comprehensive Studies of Solar Activity on the CORONAS-F Satellite"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2002ESASP.506..915Z "SPIRIT X-ray telescope/spectroheliometer results"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/2012SoPh..276...97R "Investigation of Hot X-Ray Points (HXPs) Using Spectroheliograph Mg xii Experiment Data from CORONAS-F/SPIRIT"]<br />
<br />
[5] [http://adsabs.harvard.edu/abs/2014ApJ...784...49C "Active Region Emission Measure Distributions and Implications for Nanoflare Heating"]<br />
<br />
[6] [http://adsabs.harvard.edu/cgi-bin/bib_query?arXiv:1810.04952 "Estimate of the Upper Limit on Hot Plasma Differential Emission Measure (DEM) in Non-Flaring Active Regions and Nanoflare Frequency Based on the Mg xii Spectroheliograph Data from CORONAS-F/SPIRIT"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Remembering_Marcos_Machado_via_his_researchRemembering Marcos Machado via his research2018-12-07T20:14:10Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Remembering Marcos Machado via his research<br />
|number = 336<br />
|first_author = Hugh Hudson <br />
|second_author = <br />
|publish_date = 12 November 2018<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::337]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::335]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
Marcos Machado, an old friend and colleague, passed away in September,<br />
2018, much to our sorrow.<br />
Please see the further detail in the<br />
[http://spd.stanford.edu//SolarNews/2018/20181001.html#section_mandrini obituary]<br />
written by his first PhD student, <br />
[https://www.iau.org/administration/membership/individual/1216/ Cristina Mandrini].<br />
<br />
Marcos and I shared common interests but had very different backgrounds<br />
- he, as an astronomer, actually understood radiative-transfer theory.<br />
Thus I (best defined then as a cosmic-ray physicist) learned a great<br />
deal about how solar flares work from him, and his great work from<br />
his first active period as a scientist included a really good stab<br />
at defining the solar atmosphere during the impulsive phase of a<br />
flare (Ref. [1]). Later great works involved solar X-ray astronomy, <br />
and then (to our regret) he moved into administrative work.<br />
<br />
We are very fortunate that at the end he again had enthusiastically<br />
embraced flare research, specifically re-visiting his early work<br />
in EUV astronomy - but now using the very precise <br />
[http://lasp.colorado.edu/home/missions-projects/quick-facts-sdo-eve/ SDO/EVE] data,<br />
which give Sun-as-a-star views of the <br />
[https://en.wikipedia.org/wiki/Extreme_ultraviolet EUV], a most crucial spectral<br />
band since it shows us the behavior of the <br />
[https://en.wikipedia.org/wiki/Solar_transition_region transition region] of<br />
the solar atmosphere, though without imaging. <br />
<br />
Ref. [2] captures this work, published posthumously thanks to Ryan <br />
Milligan and Paulo Sim&otilde;es.<br />
It comes just 40 years after Ref. [3], a work of major significance on<br />
the very same topic, published during Marcos's postdoctoral stay<br />
at Harvard.<br />
Note that this was actually the same institution in which <br />
[https://en.wikipedia.org/wiki/Theodore_Lyman Lyman] identified<br />
these features early in the 20th century.<br />
<br />
== The Lyman Continuum ==<br />
<br />
The new work examines the behavior of the <br />
[https://en.wikipedia.org/wiki/Lyman_continuum_photons Lyman continuum] in a <br />
representative set of strong solar flares.<br />
As the basic continuum emission (recombination to the ground state) of the most<br />
abundant element (in the nearest star) this spectral feature cries<br />
out for deep and detailed understanding.<br />
Figure 1 shows some EVE data.<br />
In thermal equilibrium, the population of the excited level <br />
should reflect the kinetic temperature of the electrons in the<br />
plasma, and indeed its slope approximates a blackbody<br />
[https://en.wikipedia.org/wiki/Max_Planck Planck] function<br />
that reflects the physical temperature in the region around<br />
[optical depth] unity.<br />
<br />
[[File:336f0.png|400px|thumb|center|Figure 1. Upper, time profiles<br />
of the flare SOL2017-09-06, a recent very powerful event.<br />
Lower, the EVE spectra at two times, showing the distinct presence<br />
of the Lyman continuum shortward of 912 &Acirc; both in the quiet<br />
Sun (black), and strongly enhanced during the flare (red).<br />
]]<br />
<br />
Figure 2 contrasts the Lyman continuum spectral parameters for the gradual<br />
and impulsive phases of SOL2017-09-06.<br />
The extremely high color temperature in the impulsive-phase spectrum<br />
(right) probably sets a record.<br />
Is it because the continuum now forms above the height of hydrogen<br />
temperature regulation, or is it a reflection of non-thermal physics?<br />
<br />
[[File:336f2.png|500px|thumb|center|Figure 2. <br />
Lyman continuum contrasted between the gradual phase (left) and impulsive<br />
phase (right) for SOL2017-09-06, along with fit parameters.<br />
The color temperature of this major flare's impulsive phase is remarkably high.<br />
]]<br />
<br />
As shown in Ref. [4], the continuum irradiance detected by EVE should<br />
vary as <br />
[[File:336f1.png|300px|thumb|center<br />
]]<br />
relating the EVE irradiance F to the Planck function, as modified by<br />
the b<sub>1</sub> "departure coefficient" and the angular subtense of<br />
the flare.<br />
The departure coefficient here is the ratio of the population of the<br />
first excited state of hydrogen, normalized to the value expected in <br />
thermal equilibrium.<br />
But the theory assumes Maxwellian distributions throughout, of course,<br />
and in the impulsive phase of a flare we certainly could have <br />
non-equilibrium processes at work.<br />
<br />
== Conclusion ==<br />
<br />
The very nice paper (Ref. [2]) that Ryan and Paulo helped Marcos to<br />
complete really confirms some of Marcos's early work, with the<br />
very limited [https://en.wikipedia.org/wiki/Apollo_Telescope_Mount Skylab] data.<br />
It also extends it because EVE gave access to the remarkably <br />
powerful flare, SOL2017-09-06. <br />
We greatly regret the loss of a friend but also the re-awakened <br />
leadership he demonstrated in this wonderful observation.<br />
It points the way to what we could learn from the Lyman continuum <br />
if only we had imaging spectroscopy to augment EVE's temporal <br />
resolution. <br />
In the meanwhile we must rely upon a combination of intuition and<br />
modeling, still in the style Marcos brought to this subject many<br />
decades ago.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/1975SoPh...42..395M "Flare model chromospheres and photospheres"]<br />
<br />
[2] [http://adsabs.harvard.edu/cgi-bin/bib_query?arXiv:1810.10824 "Lyman continuum observations of solar flares using SDO/EVE"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1978SoPh...59..129M "Lyman continuum observations of solar flares"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/1970SoPh...15..120N "The Solar Lyman Continuum and the Structure of the Solar Chromosphere"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_25_Strikes_AgainCycle 25 Strikes Again2018-12-07T20:13:54Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Cycle 25 Strikes Again <br />
|number = 337<br />
|first_author = Kamil Bicz<br />
|second_author = <br />
|publish_date = 20 November 2018<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::338]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::336]]}}<br />
}}<br />
<br />
<br />
== Introduction ==<br />
<br />
This Nugget announces that we have seen second clear sign of <br />
[https://en.wikipedia.org/wiki/Solar_cycle_25 Cycle 25]<br />
sunspot activity. <br />
Nugget No. [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_Sunspot_from_Cycle_25_for_sure 321] <br />
had described first sunspot of Cycle 25, and on November 8 (2018), at a <br />
high solar latitude, something interesting again appeared on the Sun. <br />
A tiny sunspot with Cycle 25 magnetic orientation appeared, and we believe<br />
this to be the second occurrence of the new Cycle.<br />
It persisted for about 5 days and appeared to have been bigger than its <br />
predecessor. <br />
This is interesting because it is already a second sunspot of the <br />
new cycle even before the minimum of <br />
[https://en.wikipedia.org/wiki/Solar_cycle_24 Cycle 24] <br />
has been well established.<br />
<br />
== The second Cycle 25 sunspot ==<br />
<br />
When this Nugget was written, NOAA had not identified this new<br />
sunspot with an official active-region number, and in fact it never<br />
did before solar rotation eclipsed it.<br />
We are unaware of any other public announcement, so this may be another <br />
RHESSI Nugget first!<br />
But it was duly entered in Jan Ulvestadt's <br />
[http://www.solen.info/solar/cycle25_spots.html SOLEN] list, which also<br />
noted it to be the largest Cycle 25 spot yet.<br />
The magnetic polarity of this region was unmistakeably reversed, and its <br />
appearance at high latitude strongly identifies it as a piece of the new <br />
cycle. <br />
<br />
Figure 1 here shows the sunspot on October 9, 2018 and the magnetic<br />
field polarities correlated with this sunspot. <br />
<br />
[[File:337f1.png|600px|thumb|center|Figure 1: File<br />
images from the HMI intrument on SDO: left, the telltale magnetic<br />
field; right, the continuum intensity. <br />
]]<br />
<br />
At the bottom of Figure 2 we can see a reference active region,<br />
NOAA No. 12726, which appeared a few days after region described in this <br />
Nugget. <br />
Region 12726 is part of Cycle 24 because black polarity is on the right <br />
which is opposite to polarity shown on Fig. 1. <br />
The sunspot on the continuum image can be hard to see but referring it to <br />
the location of magnetic fields can help; it's really there.<br />
<br />
[[File:337f2.png|500px|thumb|center|Figure 2: File<br />
images from the <br />
[http://hmi.stanford.edu HMI] intrument on <br />
[https://sdo.gsfc.nasa.gov SDO]. <br />
]]<br />
<br />
Sadly, nothing happened... the new region did not gave us a flare. <br />
The region simply appeared, and then the sunspot fragmented and disappeared<br />
on the disk in a few days.<br />
But it gave us hope for new regions (bigger, and with flares), hopefully <br />
in the close future.<br />
<br />
== Conclusion ==<br />
<br />
This sunspot was properly tabulated in the meticulous<br />
[http://www.solen.info/solar/ SOLEN] page of Jan Alvestad. <br />
We thank him for his diligent monitoring of solar activity<br />
As a closing thought... does this make a new <br />
[http://www.solarstorms.org/SunLikeStars.html Maunder Minimum] seem<br />
less likely?</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Neutron_Production_in_Solar_FlaresNeutron Production in Solar Flares2018-12-07T20:13:33Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Neutron Production in Solar Flares <br />
|number = 338<br />
|first_author = Ron Murphy <br />
|second_author = Gerry Share <br />
|publish_date = 27 November 2018<br />
|next_nugget ={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::339]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::337]]}}<br />
}} <br />
<br />
== Introduction ==<br />
<br />
Measurements of [https://simple.wikipedia.org/wiki/Solar_flare solar-flare] <br />
[https://simple.wikipedia.org/wiki/Neutron neutrons], along with measurements of<br />
other flare emissions, are important diagnostic tools for understanding<br />
the flare process in general and ion acceleration in particular.<br />
There have been several RHESSI Nuggets written about detections of<br />
neutrons from solar flares<br />
(e.g. [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Highly_significant_detection_of_solar_neutrons A],<br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Solar_Cosmic_Rays,_Neutrons,_and_Fermi_Gamma-Rays B],<br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Solar_flare_neutrons_observed_on_the_ground_and_in_space C]).<br />
In this Nugget we discuss the physics<br />
of neutron production, and describe comparison of neutron calculations<br />
with measurements can be used to learn about ion acceleration in<br />
solar flares.<br />
<br />
In the standard solar-flare model (Ref. [1]), electrons<br />
and ions are accelerated somehow via<br />
[magnetic reconnection], probably<br />
near the tops of closed coronal magnetic loops. <br />
These particles travel down<br />
the loop legs and interact with material at the loop footpoints.<br />
Nuclear interactions of the ions then produce excited and radioactive<br />
nuclei, neutrons, and pions.<br />
<br />
== Neutron Production ==<br />
<br />
Neutrons are produced at mostly chromospheric densities (similar<br />
to excited-nuclei production) by numerous reactions of accelerated<br />
protons, <sup>3</sup>He, and alpha particles (<sup>4</sup>He)<br />
with ambient H, He and heavier elements. <br />
For typical solar abundances and ion energy spectra,<br />
hundred-MeV neutrons are produced mostly by <br />
(<sup>1</sup>H,<sup>4</sup>He), (<sup>4</sup>He,<sup>1</sup>H), and<br />
(<sup>4</sup>He,<sup>4</sup>He) reactions.<br />
At higher neutron energies, proton-proton (<sup>1</sup>H,<sup>1</sup>H)<br />
reactions contribute significantly, while at lower energies, alpha<br />
reactions with heavy elements are important. <br />
Being neutral, neutrons travel in straight lines independent of the <br />
magnetic field. <br />
Free of the nucleus, a neutron decays radioactively with a mean lifetime of <br />
about 882 s, into a proton, an electron and a neutrino.<br />
<br />
[[File:338f1.png|600px|thumb|center|Figure 1: <br />
Neutron production in solar flares. IMF shows the typical spiral pattern<br />
of the interplanetary magnetic field, which guides charged particles<br />
but not neutrons (or photons).<br />
]]<br />
<br />
Solar-flare neutrons can be detected both and directly and indirectly <br />
(see Figure 1). <br />
Some of the higher-energy (>30 MeV) neutrons that have initial directions <br />
away from the Sun can survive the transit to Earth and be directly <br />
detected with instruments in orbit and, for flares producing sufficient <br />
high-energy (>200 MeV) neutrons, with neutron monitors on the<br />
ground. <br />
Neutrons with energies >30 MeV have been observed at 1 AU from several<br />
large flares since the initial discovery (Ref. [2]).<br />
Recently, solar-flare neutrons have also been detected<br />
by a the neutron telescope onboard the International Space Station; see an [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Solar_Cosmic_Rays,_Neutrons,_and_Fermi_Gamma-Rays earlier Nugget].<br />
For typical solar-flare ion spectra, the energies of the ions <br />
producing neutrons detectable at Earth are generally greater <br />
than about 100 MeV/nucleon<br />
<br />
Lower-energy neutrons do not survive the transit due to decay and<br />
can only be directly detected with instruments in the inner<br />
heliosphere. <br />
For example, neutron-decay products can be detected by charged-particle<br />
detectors in space (e.g., Ref. [3]).<br />
<br />
Neutrons having initial directions into the Sun may be indirectly<br />
detected because some of them can be captured on ambient hydrogen<br />
in the solar atmosphere, producing deuterium with the binding energy<br />
appearing as the strong gamma-ray line at 2.223 MeV. <br />
This line was among the first gamma-ray lines observed from a <br />
solar flare (Ref. [4]).<br />
Capture occurs after the neutrons have thermalized via scattering<br />
deep in the solar atmosphere (photosphere) where the high density<br />
favors efficient capture over decay. <br />
Although the neutron-capture line is intrinsically one of the strongest <br />
lines produced in flares, the deep capture results in significant <br />
attenuation of the line, particularly for flares occurring near the <br />
solar limb where the path-length to Earth is greatest. <br />
This is in contrast to nuclear deexcitation lines (such as <br />
<sup>14</sup>C 4.44 MeV), which are produced at chromospheric densities <br />
resulting in little attenuation of the lines unless the flare is <br />
beyond the solar limb. <br />
The ion energies responsible for the neutrons producing the<br />
capture line are typically 10-50 MeV nucleon1. <br />
Higher-energy ions produce neutrons that are captured deeper in the <br />
photosphere, resulting in stronger attenuation, and so do not contribute<br />
significantly to the observed capture line. <br />
Lower-energy neutrons tend to decay rather than be captured.<br />
<br />
== Neutron Diagnostics ==<br />
<br />
The shape of the neutron kinetic-energy spectrum depends most<br />
significantly on the shape of the accelerated-ion energy spectrum.<br />
Figure 2 (left) shows calculated neutron energy spectra escap- ing the Sun<br />
(before losses due to neutron decay) from interactions of accelerated<br />
ions directed downward into the Sun having assumed power-law<br />
kinetic-energy spectral indexes of 3 and 5 at two flare locations<br />
on the solar disk as viewed from Earth: disk center and limb. <br />
While low-energy neutrons tend to be produced isotropically, higher energy<br />
neutrons tend to be emitted more in the same direction as the<br />
interacting ion, with this tendency increasing with increasing ion<br />
energy. <br />
As a result, escaping neutron spectra from disk-centered flares tend to be <br />
steeper than those from limb flares.<br />
<br />
[[File:338f23.png|600px|thumb|center|Figure 2: <br />
Left, neutron energy spectra escaping from the Sun before losses due to <br />
neutron decay) for two accelerated-ion power-law spectral indexes s and <br />
two flare locations: disk center (&theta;<sub>obs</sub> = 0<sup>&circ;</sup>) <br />
and limb (&theta;<sub>obs</sub> = 85<sup>&circ;</sup>). <br />
Right, Ratio of the neutron-capture line and <sup>12</sup>C 4.44 MeV line <br />
fluxes for several flare locations as a function of the accelerated-ion power<br />
law spectral indexes s.<br />
]]<br />
<br />
Comparing such calculated neutron energy spectra with measured<br />
spectra (corrected for neutron decay) provides a direct means for<br />
determining both the accelerated-ion energy spectrum at energies<br />
greater than about 50 MeV/nucleon and the number of accelerated ions<br />
required to produce the neutrons.<br />
When no spectral measurements are available, neutron-arrival time histories <br />
can still be used to obtain some spectral information: if the production<br />
duration (as implied by a measured hard X-ray or gamma-ray flux<br />
time history) is known to be short, the arrival times of the neutrons<br />
can be used to infer their energies via time-of-flight.<br />
<br />
The energies of the ions producing the neutrons responsible for the <br />
neutron-capture line are higher than those producing typical <br />
deexcitation lines (5-20 MeV/nucleon). <br />
The flux ratio of the neutron-capture line and a<br />
deexcitation line is therefore sensitive to the shape of the<br />
accelerated-ion energy spectrum across the 5-50 MeV/nucleon range.<br />
Shown in Figure 2 (right) is the calculated ratio of the neutron-capture<br />
line flux to the <sup>12</sup>C 4.44 MeV deexcitation line <br />
flux for several flare locations on the Sun as a function of the index of the<br />
accelerated-ion spectrum, assumed here to be a power law. <br />
These curves can be used with measured flux ratios to determine the ion<br />
spectral index across the 5-50 MeV/nucleon energy range. <br />
Comparing calculated neutron-capture fluxes with an observed flux then<br />
determines the required number of accelerated ions.<br />
<br />
Ion spectral indexes derived from this ratio can be<br />
compared with indexes derived for different ion-energy ranges,<br />
providing ion spectral information across a wider energy range. <br />
For example, the flux ratio of the <sup>20</sup>Ne 1.63 MeV and <br />
<sup>16</sup>O 6.13 MeV deexcitation lines is sensitive to the ion <br />
spectral shape across the 2-20 MeV/nucleon energy range. <br />
This ratio is, of course, strongly dependent on the ambient <br />
(chromospheric) [Ne]/[O] abundance ratio. <br />
The [https://en.wikipedia.org/wiki/Solar_Maximum_Mission SMM] <br />
made measurements of this flux ratio from and the <br />
neutron-capture/4.44 MeV line ratio. <br />
These showed that the ion spectrum could be an unbroken power law across<br />
the full 2-50 MeV/nucleon energy range only if the ambient [Ne]/[O]<br />
abundance ratio is 0.25, significantly more than the accepted value<br />
of 0.15. <br />
Note that the chromospheric-photospheric [Ne]/[O] abundance ratio<br />
cannot be directly measured by other means. <br />
If the ambient [Ne]/[O] abundance is instead 0.15, the ion spectrum <br />
below ~20 MeV/nucleon must steepen, which would have a significant <br />
impact on the number of and the energy contained in low-energy <br />
ions accelerated in solar flares.<br />
<br />
The neutron-capture line flux can also be compared with the pion-decay<br />
flux to determine the accelerated-ion spectrum shape at higher<br />
energies. <br />
In solar flares, pions are produced primarily by proton-proton<br />
and proton-alpha (and alpha-proton) interactions. <br />
The threshold energy for pion production is ~200 MeV for <br />
proton-alpha interactions and ~300 MeV for proton-proton interactions. <br />
The ratio of the neutron-capture line flux and the >100 MeV pion-decay <br />
flux is therefore sensitive to the accelerated proton spectral shape <br />
across the ~30-300 MeV/nucleon energy range. <br />
Figure 3 (left) shows this ratio calculated for several flare<br />
locations on the Sun as a function of the ion power-law spectral<br />
index. <br />
These curves can be used with measured flux ratios to determine<br />
the ion spectral index across the 30-300 MeV/nucleon energy range. <br />
The accelerated-proton spectrum above 300 MeV/nucleon can be determined by<br />
directly fitting measured gamma-ray spectra with calculations of<br />
the pion-decay spectrum whose shape is sensitive to the proton<br />
spectral shape. <br />
Combining deexcitation-line, neutron-capture line<br />
and pion-decay emission measurements allows the shape of the<br />
accelerated-ion energy spectrum to be determined across the broad<br />
energy range of 2 to >300 MeV/nucleon. <br />
<br />
[[File:338f45.png|600px|thumb|center|Figure 3: <br />
Left, ratio of the >100 MeV pion-decay and neutron-capture<br />
line fluxes for several flare locations as a function of the<br />
accelerated-ion power law spectral index s.<br />
Right, calculated photon spectra at Earth from 2.223 MeV neutron-capture<br />
gamma rays for several flare observation angles (&\theta;<sub>obs</sub>). <br />
]]<br />
<br />
Because the neutrons thermalize before being captured, the<br />
neutron-capture line shape only reflects thermal broadening due to<br />
the photospheric temperature and is very narrow (<100 eV). <br />
The shape has no diagnostic value, in contrast to deexcitation-line shapes<br />
which depend on a number of flare parameters such as the accelerated-ion<br />
spectral shape and interacting angular distribution, and the<br />
accelerated alpha/proton ratio.<br />
<br />
== Compton-Scattering of the Neutron-Capture Line ==<br />
<br />
Because neutron capture occurs deep in the solar atmosphere, both<br />
downward- and upward-directed photons can Compton-scatter once<br />
or several times before escaping. <br />
This attenuates the line flux and produces a continuum of photons at <br />
energies below the line energy.<br />
(This is in contrast to the de-excitation lines, which are produced<br />
higher in the atmosphere, and Compton scattering of only the<br />
downward-directed photons is significant.) <br />
Figure 3 (right) shows calculated photon spectra escaping the Sun <br />
from scattering of 2.223 MeV neutron-capture line gamma rays <br />
for several flare locations on the solar disk. <br />
This Compton continuum may be diagnostically useful in<br />
that it may still be detectable for flares beyond the solar limb<br />
while the capture line itself is not.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2008LRSP....5....1B "Flare Observations"]<br />
<br />
[2] [http://cdsads.u-strasbg.fr/abs/1982ApJ...263L..95C "A direct observation of solar neutrons following the 0118 UT flare on 1980 June 21"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1966JGR....71.1305R "Effect of the Interplanetary Magnetic Field on Solar Neutron-Decay Protons"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/1973NASSP.342..285C "Solar Gamma Ray and Neutron Observations"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/The_flight_of_FOXSI-3The flight of FOXSI-32018-12-07T20:12:37Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = The flight of FOXSI-3: Single photon counting and direct focusing across hard and soft energies <br />
|number = 340<br />
|first_author = Lindsay Glesener <br />
|second_author = Noriyuki Narukage<br />
|publish_date = 10 December 2018<br />
|next_nugget = TBD<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::339]]}}<br />
}} <br />
<br />
== Introduction ==<br />
<br />
The FOXSI-3 <br />
[https://en.wikipedia.org/wiki/Sounding_rocket sounding rocket] <br />
flew for a successful 6-minute observation of the Sun on September 7, 2018 <br />
from the <br />
[https://www.nps.gov/whsa/learn/historyculture/white-sands-missile-range.htm White Sands Missile Range].<br />
FOXSI stands for the Focusing Optics X-ray Solar Imager, and it is the<br />
first experiment to optimize direct-focusing optics for the Sun in<br />
hard X-rays (HXRs). <br />
This technology offers dramatically increased<br />
sensitivity, allowing FOXSI to investigate coronal heating in active<br />
regions and to search for non-thermal signatures in the quiet Sun.<br />
When eventually implemented on a space mission, FOXSI will investigate<br />
particle acceleration in solar flares with unprecedented sensitivity.<br />
<br />
The FOXSI rocket experiment has had two prior successful flights already,<br />
in 2012 and 2014, resulting in the first focused HXR image of the<br />
Sun (Ref. [1]) and the most direct measurement to date of nanoflare-heated<br />
plasma in an active region (Ref. [2]).<br />
New for the third payload were upgrades to the optics and detectors, <br />
an improved white-light camera for better alignment, collimators to <br />
block stray light (ghost rays), and a focusing soft X-ray <br />
telescope that measures individual photons in the 0.5-5 keV range. <br />
While telescopes of this type have been used in astrophysics for decades <br />
(i.e. [http://chandra.harvard.edu Chandra] or <br />
[http://sci.esa.int/xmm-newton/ XMM-Newton]), <br />
the bright soft X-ray fluxes of the Sun long prevented their solar use. <br />
FOXSI's PhoEnIX (Photon Energy Imager in soft X-rays) instrument solves this <br />
difficulty with a <br />
[https://hinode.nao.ac.jp/en/news/topics/foxsi-3-180907/ high-speed CMOS X-ray camera].<br />
<br />
== FOXSI-3 Observations ==<br />
<br />
The Sun on September 7 was mostly quiet, exhibiting <br />
[https://en.wikipedia.org/wiki/Coronal_hole coronal holes],<br />
[http://www.oxfordreference.com/view/10.1093/oi/authority.20110803095640219i bright EUV points], <br />
and a 4-month old <br />
[http://www.scholarpedia.org/article/Solar_activity active region] late in <br />
its lifetime (and no longer numbered by<br />
[https://www.swpc.noaa.gov NOAA]). <br />
No flares were observed. <br />
Due to the low level of activity, FOXSI's six HXR telescopes registered<br />
only a few photons. <br />
The soft X-ray <br />
[https://hinode.nao.ac.jp/en/news/topics/foxsi-3-180907 PhoEnIX]<br />
telescope though, operating at lower energies, produced detailed images of the <br />
Sun, as shown in Figure 1. <br />
These are spectroscopic images produced from single<br />
photons, meaning that time profiles and energy spectra are available<br />
for every pixel in the image (see Figure 2). <br />
Figure 1 is a mosaic of three different targets. <br />
In addition to demonstrating the high-quality performance of PhoEnIX, <br />
it also demonstrates a striking lack of ghost ray background due to the <br />
successful implementation of FOXSI-3's collimators (Ref. [3]).<br />
<br />
[[File:340f1.png|400px|thumb|center|Figure 1: <br />
FOXSI/PhoEnIX soft X-ray image produced from a mosaic of three<br />
different targets. <br />
This is the first soft X-ray image of the Sun<br />
produced via single-photon counting.<br />
]]<br />
<br />
[[File:340f2.png|700px|thumb|center|Figure 2: FOXSI observes the<br />
Sun via single photon counting in both hard and soft X-rays. <br />
These panels show data from the PhoEnIX soft X-ray telescope. <br />
The top panels show snapshots of single-photon hits on the fast CMOS detector. These photons can then be binned in space, time, and energy to produce<br />
images, time profiles, and spectra (bottom row).<br />
]]<br />
<br />
== FOXSI-3 Science ==<br />
<br />
Analysis of the data acquired during FOXSI-3's six-minute observation is <br />
currently ongoing. <br />
In addition to FOXSI's own hard and soft X-ray data, several coordinated <br />
instruments observed the primary target, including <br />
[http://global.jaxa.jp/projects/sat/solar_b/ Hinode]/EIS and XRT, <br />
[http://iris.lmsal.com IRIS], and <br />
[https://www.nustar.caltech.edu NuSTAR]. <br />
With this wealth of multiwavelength data, we expect results in the following <br />
areas:<br />
<br />
&bull; Determination of the temperature structure and heating<br />
mechanisms for an aged active region.<br />
<br />
&bull; Measurement of temperature structure of bright points<br />
observed in EUV and X-rays.<br />
<br />
&bull; Limits on non-thermal emission from the quiet Sun, with<br />
corresponding consequences for how the quiet corona is heated.<br />
<br />
== Conclusions ==<br />
<br />
The [http://foxsi.umn.edu FOXSI-3] sounding rocket is a collaboration between <br />
the University of Minnesota, University of California Berkeley, NASA/Marshall,<br />
NASA/Goddard, University of Tokyo/Kavli IPMU, Nagoya University,<br />
Tokyo University of Science, JAXA/ISAS, and the National Astronomical<br />
Observatory of Japan. <br />
With three successful flights demonstrating the feasibility and<br />
scientific promise of focusing optics for high-energy solar<br />
observation, FOXSI is currently competing in a concept study for a<br />
Small Explorer spacecraft led by NASA/GSFC.<br />
<br />
The FOXSI-3 sounding rocket is funded by NASA LCAS grant NNX16AL60G.<br />
We acknowledge the NSF for its support of space physics at UMN via<br />
a faculty development grant AGS-1429512. The FOXSI team is grateful<br />
to the NSROC teams at WSMR and Wallops for the excellent operation<br />
of their systems. Furthermore, the authors would like to acknowledge<br />
the contributions of each member of the FOXSI experiment team to<br />
the project, particularly our team members at ISAS and Kavli IPMU<br />
for the provision of Si and CdTe detectors and at MSFC for the<br />
fabrication of the focusing optics. This work was supported by<br />
JSPS KAKENHI Grant Numbers JP17H04832, JP16H02170, JP16H03966,<br />
JP24244021, JP20244017, and World Premier International Research<br />
Center Initiative (WPI), MEXT, Japan. PhoEnIX work was supported<br />
by JSPS KAKENHI Grant Numbers JP18H03724, JP18H05463, JP17H04832,<br />
JP16H02170, JP15H03647, JP21540251.<br />
<br />
== References ==<br />
<br />
[1] [https://doi.org/10.1088/2041-8205/793/2/L32 "First images from the Focusing Optics X-ray Solar Imager"]<br />
<br />
[2] [https://doi.org/10.1038/s41550-017-0269-z "Detection of nanoflare-heated plasma in the solar corona by the FOXSI-2 sounding rocket"]<br />
<br />
[3] [https://doi.org/10.1117/12.2274675 "Methods for reducing singly reflected rays on the Wolter-I focusing mirrors of the FOXSI rocket experiment"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Stellar_Flares_and_StarspotsStellar Flares and Starspots2018-12-07T20:11:58Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Stellar Flares and Starspots<br />
|number = 339<br />
|first_author = Lauren Doyle <br />
|second_author = <br />
|publish_date = 3 December 2018<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::340]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::338]]}}<br />
}} <br />
<br />
== Introduction ==<br />
<br />
Solar flares are well-known phenomena and have been studied<br />
extensively for the past 150 years in many ways, for example by RHESSI. <br />
We now know there is a direct connection between the Sun's magnetic field <br />
and [http://www.scholarpedia.org/article/Solar_activity solar activity].<br />
Similarly, flares have been seen from very low mass stars for many<br />
decades with extensive observational evidence that magnetic fields<br />
drive them (e.g., Ref. [1]). <br />
These cool, small, main sequence stars become fully<br />
convective for masses lower than ~0.3 M<sub>&#9737;</sub>. <br />
However, despite the<br />
lack of a <br />
[https://en.wikipedia.org/wiki/Tachocline tachocline] we still observe <br />
strong flaring activity from these stars.<br />
Solar magnetism probably comes from an internal <br />
[https://en.wikipedia.org/wiki/Tachocline dynamo] <br />
action, thought to involve the strong velocity shears found at the<br />
tachocline.<br />
Thus the presence of flares on a fully convective star - which would have<br />
no tachocline - suggests that the magnetic field <br />
must be generated through a different mechanism compared with the Sun.<br />
<br />
Through 10 years of observations, the <br />
[https://en.wikipedia.org/wiki/Kepler_(spacecraft) Kepler]<br />
observatory has revolutionised the field of <br />
[https://spaceplace.nasa.gov/all-about-exoplanets/en/ exoplanet] research <br />
and the accompanying stellar physics. <br />
For four years it focussed<br />
on the same patch of sky, providing data for over 100,000 stars. <br />
In 2014 it lost the use of two reaction wheels and was repurposed as<br />
"K2," changing the way in which it observed. <br />
Now four further years later, it has observed 18 new star fields along the <br />
ecliptic, providing photometric data for thousands of low mass stars. <br />
Amongst several of the K2 fields, we had a helping hand in putting these <br />
low mass stars on the observing programme.<br />
<br />
Many of the photometric lightcurves of low mass M dwarf stars from<br />
K2 possess substantial variations in brightness due to the presence<br />
of a large, dominant [https://en.wikipedia.org/wiki/Starspot starspot] <br />
on the stellar surface. <br />
As the starspot moves in and out of view while the star rotates, <br />
it produces nearly sinusoidal changes in brightness over time; see Figure 1. <br />
In our study, we used a sample of low mass stars observed by K2 to <br />
investigate their flare properties. <br />
The vast majority of lightcurves made using Kepler<br />
and K2 had one photometric point only every 30 minutes. For a small<br />
number of stars, including our sample, we obtain one point every<br />
minute, which is essential to identify low-amplitude short-duration<br />
flares.<br />
<br />
[[File:339f1.png|400px|thumb|center|Figure 1: <br />
Sketch showing the location of the dominant starspot on a low-mass star,<br />
and the corresponding photometric lightcurve. <br />
This highlights the observed change in brightness as the cooler and <br />
darker starspot passes across the stellar disk.<br />
]]<br />
<br />
== Flares from Low Mass Stars ==<br />
<br />
Identifying flares within the lightcurve<br />
can be a time-consuming task and so, we utilise "Flares By Eye (FBEYE)",<br />
an IDL suite of programs written by J. R. A. Davenport (Ref. [2]).<br />
From this, we can derive the duration and phase of the flares as well <br />
as the energy from ground-based multi-epoch<br />
surveys such as that of [https://panstarrs.stsci.edu PanStarrs]. <br />
We find there is a decline in flaring activity after rotation periods<br />
longer than about 10 days, which is consistent with previous work (Ref. [3]). <br />
However, the relationship between spottedness and a star's activity is <br />
complex and depends on other factors such as age and mass. <br />
This has been explored by many others using K2 observations of stellar <br />
clusters of varying ages.<br />
<br />
== A Surprise Finding ==<br />
<br />
In solar physics, the relationship between sunspots and solar flaring<br />
activity has been studied for decades and it is generally accepted<br />
that they are closely related, but physically very different phenomena.<br />
In low mass stars, we observe periodic modulations in their brightness <br />
caused by starspots as the star rotates. <br />
In Figure 2 we show an example of this and also observe<br />
flares occurring continuously in the lightcurve, even at a peak in<br />
the rotational modulation. <br />
If we assume the physical processes which occur in solar flares are to hold <br />
for stellar flares, then one might expect to see the flares originating from <br />
active regions which typically host spots. <br />
Consequently, there should be a correlation<br />
between the rotational phase (i.e. when the spot is most visible)<br />
and the number of flares. <br />
<br />
[[File:339f2.png|700px|thumb|center|Figure 2: <br />
Section of a K2 short-cadence lightcurve for the star <br />
[http://simbad.u-strasbg.fr/simbad/sim-id?Ident=LHS+3197 GJ 3954 (EPIC 205467732)] <br />
from Field 2, covering approximately seven days, with a rotation period of <br />
1.321 days. <br />
The black points represent the K2 data points and the red line is a <br />
smoothed version of the time-series data. <br />
Throughout this lightcurve there are larger flares present during the <br />
maximum peak of rotation when the starspot is least visible, and <br />
smaller flares at all rotational phases.<br />
]]<br />
<br />
In order to investigate this behaviour further, we conducted a simple <br />
&chi;<sup>2</sup> test to determine whether the<br />
phase distribution is random, as shown in Figure 3. <br />
None of the stars in our sample show any preference for rotational phase <br />
dependence, suggesting that many of the flares do not originate from nearby<br />
the large dominant spot. <br />
The question remains, what are the mechanisms responsible for causing<br />
the generation of flares in these low mass active stars?<br />
<br />
[[File:339f3.png|500px|thumb|center|Figure 3: <br />
Top panel: the normalised, phase folded, binned lightcurve of the flare star <br />
[http://simbad.u-strasbg.fr/simbad/sim-id?Ident=LHS+3197 GJ 3954]. <br />
Bottom panel: the phases of the flares with respect to their energy. <br />
The data in each plot have been plotted twice, covering the rotational phases<br />
from 0.0-2.0 where 1.0-2.0 is simply a repeat. <br />
There is no correlation between flare number and rotational phase.<br />
]]<br />
<br />
== Conclusions ==<br />
<br />
To summarise, we find no correlation between the rotational phase<br />
and the number of flares in our sample of M dwarfs observed in short<br />
cadence by K2. <br />
Interestingly, our result indicates flares do not<br />
originate from the active region hosted by the large starspot. <br />
We outline three scenarios to explain this: polar spots, stellar<br />
binarity systems and the presence of orbiting. To test the three<br />
scenarios further we would need to investigate the inclination and<br />
potential star-planet or binary stars in greater detail.<br />
<br />
Our initial study (Ref. [4]) assessed the flare characteristics of 34 <br />
low-mass stars using K2 data from observing fields 1-9.<br />
Now we are in the process of repeating this analysis for stars<br />
observed in fields 10-18. <br />
More recently, [https://tess.mit.edu TESS] was launched and will be providing<br />
exciting new photometric data.<br />
Data becoming available in early 2019 will allow us to expand<br />
our sample of low mass stars. <br />
However, we will extend this research to solar-like stars, which will provide <br />
an even closer comparison to our Sun.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2014ApJ...797..121H "Kepler Flares. I. Active and Inactive M Dwarfs"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2014ApJ...797..122D "Kepler Flares. II. The Temporal Morphology of White-light Flares on GJ 1243"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2016MNRAS.463.1844S "A path towards understanding the rotation-activity relation of M dwarfs with K2 mission, X-ray and UV data"]<br />
<br />
[4][http://adsabs.harvard.edu/abs/2018MNRAS.480.2153D "Investigating the rotational phase of stellar flares on M dwarfs using K2 short cadence data"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Photospheric_response_to_a_flarePhotospheric response to a flare2018-10-10T20:23:06Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Photospheric response to a flare <br />
|number = 332<br />
|first_author = Mike Wheatland <br />
|second_author = <br />
|publish_date = 17 September 2018<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::333]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::331]]}}<br />
}}<br />
<br />
<br />
== Introduction ==<br />
<br />
[https://en.wikipedia.org/wiki/Solar_flare Solar flares] <br />
involve sudden conversion of magnetic energy into other<br />
forms in the solar corona. <br />
In particular, flares accelerate large numbers of electrons to 10-100 keV, <br />
producing hard-X-ray emission in the <br />
[https://solarscience.msfc.nasa.gov/chromos.shtml low solar atmosphere]. <br />
Flares introduce sudden and permanent changes in the <br />
[https://en.wikipedia.org/wiki/Stellar_magnetic_field photospheric magnetic field]. <br />
Vector magnetogram observations show that the predominant change is in the <br />
horizontal magnetic field (parallel to the photosphere), which tends <br />
to increase along a flaring neutral line. <br />
By definition the neutral line marks the change of polarity of the field,<br />
dividing outward from inward domains.<br />
Flares can also produce sudden photospheric motion, strikingly illustrated by <br />
the rotation of a sunspot in response to a flare (Refs. [1,2])<br />
an event which was likened to "the tail wagging the dog" (Ref. [3]). <br />
The sunspot was observed to differentially rotate as the flare ribbons <br />
swept across it. <br />
Figures 1 and 2 show photospheric data for this event.<br />
<br />
[[File:332f1.png|600px|thumb|center|Figure 1: Changes in the vertical<br />
component of the magnetic field (left) and the vertical electric<br />
current density (right) in the flare SOL2015-06-22, using<br />
[http://hmi.stanford.edu HMI] data from the<br />
[https://sdo.gsfc.nasa.gov SDO] space observatory. <br />
The rotating sunspot reported by Ref. [1] is in the box. <br />
The electric current density is observed to suddenly increase close to<br />
the neutral line (black), coincident with the flare. <br />
]]<br />
<br />
[[File:332f2.jpg|700px|thumb|center|Figure 2: <br />
The vector change in the horizontal magnetic field for the SDO/HMI data <br />
shown in Figure 1. <br />
The horizontal field increases over a broad region along the neutral line. <br />
The largest changes are about 1000 gauss.<br />
]]<br />
<br />
The photospheric field changes are interpreted as the photospheric response <br />
to coronal magnetic restructuring due to <br />
[https://en.wikipedia.org/wiki/Magnetic_reconnection magnetic reconnection]. <br />
Here we consider a simple 2D model for the changes in the field and flows <br />
introduced at the photosphere, in terms of a large-amplitude shear Alfvén<br />
wave impacting the photosphere (Ref. [4]). <br />
<br />
== Large amplitude shear Alfvén wave model ==<br />
<br />
We consider a simple 2D model in which the photosphere is the z = 0 plane, <br />
and the x-axis is directed away from the neutral line. <br />
Figure 3 illustrates the model. <br />
A large-amplitude shear Alfvén wave is assumed to propagate downwards<br />
at the coronal Alfvén speed <sub>A1</sub>, introducing shear field and <br />
flow components B<sub>y</sub> = B<sub>1</sub> and <br />
v<sub>y</sub> = v<sub>1</sub>, respectively, behind a front. <br />
The wave front is assumed to be obliquely oriented to the photosphere, <br />
so that the changes are introduced first close to the neutral line, and <br />
then further away. <br />
The wave is partially reflected and partially transmitted at the <br />
photospheric boundary. <br />
The wave propagates at a lower Alfvén speed<br />
v<sub>A2</sub> in the sub-photosphere. <br />
Behind the reflected and transmitted fronts, there are new shear components <br />
B<sub>y</sub> = B<sub>2</sub> and v<sub>y</sub> = v<sub>2</sub>.<br />
<br />
The MHD equations and continuity imply the relationships:<br />
<br />
[[File:eq.png|400px|thumb|center|<br />
]]<br />
<br />
so that in the limit of a very dense photosphere, B<sub>2</sub> ➝ 2B<sub>1</sub> and v<sub>2</sub> ➝ 0.<br />
<br />
[[File:332f3.png|600px|thumb|center|Figure 3: <br />
The geometry of the shear Alfvén wave model. The z-direction is the local vertical direction, z = 0 is the photospheric boundary, and the x-direction is directed away from the neutral line. The point P is the location at which the wave is just impacting the photosphere. To the right of this location the shear Alfvén wave is propagating down from above, with an oblique front. To the left of this location there is a reflected and a transmitted front.<br />
]]<br />
<br />
This simple model can account for a variety of features of the observations. If the wave is incident on the photosphere on both sides of the neutral line and the change in the field is in the same direction on either side, then the plasma flow is oppositely directed on either side, reproducing the qualitative features of the observations. The shear field and flow components are related by the Walén relation v<sub>2</sub> = -v<sub>A2</sub>B<sub>2</sub>/B<sub>0</sub>. This implies changes quantitatively consistent with observed values.<br />
<br />
== Particle acceleration ==<br />
<br />
Flare ribbons are the locations of hard X-ray emission in flares. In the model, the flare ribbons are assumed to coincide with the front, and the front represents a surface current. The geometry implies a field-aligned electric field if the conductivity of the plasma is finite. <br />
Above a critical value, this field may produce electron runaway. <br />
Assuming that the dominant resistivity is due to electron-neutral collisions, we estimate that the thickness of the front must be of order 10 m to produce runaway. <br />
Although this model for particle acceleration is speculative, it has an attractive feature. <br />
Standard models for particle acceleration assume that acceleration occurs in the corona, but this raises a "number problem": the electron flux implied by the photospheric observations implies acceleration of a very large volume of electrons in the rarified corona. <br />
Acceleration in the dense lower atmosphere avoids this difficulty. <br />
However, the present model also implies an asymmetry in the direction of electron acceleration on either side of the neutral line: more hard X-rays are expected on one side of the line than the other. <br />
Figure 4 illustrates this point.<br />
<br />
[[File:332f4.JPG|700px|thumb|center|Figure 4: <br />
If the surface current density in the front (assumed to be propagating down from the corona on both sides of the neutral line, NL) produces runaway electron acceleration due to a field-aligned electric field in the atmosphere, then the electrons will be accelerated up on one side of the NL, and down on the other. <br />
More hard X-rays are expected on the down side. <br />
Which side is which depends on the shear. This figure shows a HXR producing ribbon as dark, and a non-HXR producing ribbon as white. <br />
The two possible shear directions are shown.<br />
]]<br />
<br />
== Conclusions ==<br />
<br />
Flares are observed to produce sudden changes in the photospheric field, and to introduce shear flows, along the flaring neutral line. <br />
We present a simple model for the changes in terms of a large-amplitude shear Alfvén wave impacting the photosphere. <br />
This model can account for various features of the observations. It also introduces a possible mechanism for electron acceleration in the low atmosphere, which might avoid the electron "number problem."<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2016NatCo...713104L "Flare differentially rotates sunspot on Sun's surface"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2018ApJ...853..143W "Evolution of Photospheric Flow and Magnetic Fields Associated with the 2015 June 22 M6.5 Flare"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2016NatPh..12..998A "Solar physics: When the tail wags the dog"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/2018ApJ...864..159W "Photospheric Response to a Flare"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Coronal_Hard_X-ray_Sources_RevisitedCoronal Hard X-ray Sources Revisited2018-10-10T20:22:49Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Coronal Hard X-ray Sources Revisited<br />
|number = 333<br />
|first_author = Brian Dennis<br />
|second_author = <br />
|publish_date = 24 September 2018<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::334]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::332]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
For the most part RHESSI observes hard X-rays (HXRs) from bright ''footpoint'' sources via the <br />
[http://solar.physics.montana.edu/nuggets/2000/000303/000303.html Neupert Effect].<br />
These footpoint sources lie at the intersections of coronal magnetic fields with the lower solar atmosphere.<br />
As many of these Nuggets have discussed, RHESSI also sees a variety of sources in the solar corona.<br />
<br />
A recent paper (Ref. [1]) attempts to set the record straight by reinterpreting observations of a group of<br />
flares that have been reported to have hard X-rays (HXRs) coming<br />
predominantly from the corona rather than from the more usual chromospheric<br />
footpoints (Ref. [2,3,4,5,6]). <br />
All of the 26 previously analyzed event time intervals, over 13 <br />
flares, were re-examined for consistency with a model in which electrons are accelerated near<br />
the top of a magnetic loop that has a sufficiently high density to stop most of the<br />
electrons by [https://en.wikipedia.org/wiki/Charles-Augustin_de_Coulomb Coulomb]<br />
[https://en.wikipedia.org/wiki/Coulomb_collision collisions] before they can reach the footpoints. <br />
Of particular importance in the previous analysis was the finding that the length of the coronal HXR<br />
source increased with energy in the 20 - 30 keV range. <br />
Such behavior is inconsistent with a unique thermal source in equilibrium, the size of which must generally decrease with increasing<br />
energy as the emission becomes more and more dominated by the hottest regions.<br />
The observed behavior would however be consistent with the transport of accelerated electrons through a<br />
collisional target, since higher energy electrons travel further.<br />
<br />
However, after allowing for the<br />
possibility that footpoint emission at the higher energies affects the inferred length<br />
of the coronal HXR source, and using analysis techniques that suppress the possible<br />
influence of such footpoint emission, we conclude that there is no longer evidence<br />
that the length of the HXR coronal sources increase with increasing energy. <br />
In fact, for the 6 flares and 12 time intervals that satisfied our selection criteria,<br />
the loop lengths decreased on average by 1.0 &plusmn; 0.2 arcsec between 20 and<br />
30 keV, with a standard deviation of 3.5 arcsec. <br />
We also find strong evidence that the peak of the coronal HXR source increases in altitude with increasing energy. <br />
For the thermal component of the emission, this is consistent with the standard <br />
([http://solar.physics.montana.edu/magara/Research/Topics/cshkp.html "CHSKP"])<br />
flare model in which <br />
[https://www.nasa.gov/content/goddard/science-of-magnetic-reconnection magnetic reconnection] in a coronal <br />
[https://en.wikipedia.org/wiki/Current_sheet current sheet] <br />
results in new<br />
hot loops being formed at progressively higher altitudes. <br />
The explanation for the nonthermal emission remains unclear.<br />
<br />
== Example Flare - 14/15 April 2002 ==<br />
<br />
The effect of weak footpoint emission at higher energies on the apparent source length is shown during the M3.7 flare in <br />
[https://www.solarmonitor.org/?date=20020415 NOAA AR09893] at N16W60 that peaked at 00:14 UT on 15 April 2002. <br />
The time history is shown in Figure 1 and the spectrum for the 5 min. interval at the start of the <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/%22Impulse_Response_Flares%22_and_Gamma_Rays impulsive emission] is shown in Figure 2. <br />
The count fluxes are fitted to a thermal component consisting of a differential emission measure that is a power-law in temperature (multi_therm_pow) plus a nonthermal component (thick2_vnorm) with a power-law electron spectrum. <br />
With these assumed components, the nonthermal emission dominates at energies up to ~22 keV. This is to be compared to a crossover energy of ~15 keV if an isothermal function is assumed as was done previously (Ref. [4]). Such an isothermal function is inconsistent with the change in altitude of the coronal source with energy as indicated in Figure 3, where a 12 - 25 keV image is shown with the 25 - 50 keV contours overlaid. The loop structure is clearly evident in both energy ranges but with a ~5 arcsec offset corresponding to a higher altitude in the higher energy band. Two footpoint sources appear in the 25 - 50 keV image at the expected locations with respect to the loop structures.<br />
<br />
[[File:Lc_15April2002_hsi_ql.jpg|500px|thumb|center|Figure 1:<br />
RHESSI light curves for the flare SOL2002-04-15. The color-coded curves<br />
are for the five indicated energy ranges. Counts from the front segments of all detectors<br />
except for detectors #2 and #7 were summed and divided by the summed live times and<br />
the total effective sensitive area of 35.59 cm2 per detector to give the plotted values with a<br />
4 s cadence to match the spacecraft spin period. The thin attenuators were in place above<br />
all detectors limiting the useful energy range to > 6 keV. The blue shaded areas show the<br />
first of the three time intervals between 00:00 and 00:05 UT used here and by Guo et al.<br />
(2012a,b, 2013) to determine the source dimensions. <br />
]]<br />
<br />
[[File:Sp_15april2002_000000to000500_hsi_mtpow_thick2_albedo_5to60keV_d4f.jpg|500px|thumb|center|Figure 2: <br />
RHESSI count flux spectrum for the five-minute time interval shown in Figure 1. <br />
The black histogram is the background-subtracted count flux in the front segment of Detector #4. <br />
The red histogram is the function that was fitted to the data between 6 and 50 keV. It is the sum of the following components: a multi-thermal bremsstrahlung function (green), a power-law nonthermal thick-target function (yellow), an albedo function for isotropic emission (pink), the estimated pulse pile-up contribution (purple), and two Gaussian instrumental lines (olive and brown). The background spectrum (green) determined from the nighttime period immediately prior to the are is shown with +/-1 sigma error bars. The values of all parameters used for the fit are given for each functional component.]]<br />
<br />
[[File:Im 15april2002 000000to000500 hsi memnjit 12to25to50kev d3456.jpg|500px|thumb|center|Figure 3:<br />
RHESSI 12 - 25 keV colored image overlaid with 25 - 50 keV contours for the 5-min. interval from 00:00 to 00:05 UT made using the MEM NJIT<br />
reconstruction technique with data from the front segments of detectors 3, 4, 5, and 6. Two<br />
compact footpoints are evident in the 25 - 50 keV image with the extended coronal source<br />
present at both energies but at a higher altitude in the higher energy image. ]]<br />
<br />
=== Source Length Determination ===<br />
<br />
We used the following method to estimate the length of the coronal loop source in order to minimize the effect of the footpoint emission. <br />
First, we defined an arc along the spine of the loop structure that also passed through the two footpoints. <br />
To accommodate the change in source location with energy shown in Figure 3, the arcs for each energy bin were moved to always pass through the pixel with peak emission. <br />
The resulting intensity plots along the arcs at each energy are shown in Figure 4. <br />
They were fitted with multiple Gaussians and the full width at 75%, 50%, and 20% of the peak was determined. <br />
Again, this was to minimize the effects of the footpoint emission, evident at the higher energies, on the estimate of the length of the coronal part of the source.<br />
<br />
[[File:mem_njit_event3_gaussfit_arc2.jpg|500px|thumb|center|Figure 4:<br />
Normalized X-ray intensity vs. distance measured along the loop spine that passes through the coronal source peak location at each energy<br />
Color-coded points for each energy show the intensities at each pixel along the arc in the<br />
image. The curves are the sum of the minimum number of Gaussians needed to adequately<br />
fit the data points. For each energy range, the diamonds show the location of the peak<br />
intensity; the circles and squares show the left and right location, respectively, at a given<br />
fraction of the peak intensity, with blue at 20%, orange at 50%, and red at 75%.]]<br />
<br />
The source lengths defined in this way are plotted vs. energy in Figure 5 for the three different measures of the length. <br />
The change in position of the peak is plotted in Figure 6 vs. energy. In each case, the change in length (&Delta;L) and in distance (&Delta;D) between 10 and 20 keV and between 20 and 30 keV are indicated in arcsec. <br />
For this time interval, the source length clearly decreased with increasing energy between 20 and 30 keV, contrary to the previously reported result where the measured length increased with energy. <br />
The X-ray emission in this energy range was originally thought to be nonthermal given the assumed isothermal function used for the spectral fits. <br />
The more likely multi-temperature function used for the fit shown in Figure 2 suggests that there is still considerable thermal/nonthermal overlap in this energy range with the crossover energy at ~22 keV.<br />
<br />
[[File:mem_njit_event3_fwhm_arc2.jpg|500px|thumb|center|Figure 5:<br />
Coronal source length measured along the loop spine plotted vs. photon energy. <br />
The lengths &Delta;L were determined as the distances between<br />
the positions at which the intensity along the arc dropped to 75% (blue), 50% (orange),<br />
and 20% (red) of the peak value in each energy bin. The Solid lines indicate the best fit to the data points from 10<br />
to 20 keV and from 20 to 30 keV. <br />
The changes in length (&Delta;L) over these two energy ranges<br />
are shown on the plot.]]<br />
<br />
[[File:mem_njit_event3_dist.jpg.jpg|500px|thumb|center|Figure 6:<br />
Energy dependence of the distance, &Delta;D, between the peak in coronal<br />
flux in the 20-22 keV image and the peak at other energies. <br />
The solid lines show linear fits to the data<br />
points from 10 to 20 keV and 20 to 30 keV, and the indicated values of &Delta;D<br />
are the changes in position in arcsec over the two energy ranges.<br />
]]<br />
<br />
== Results for Multiple Flares == <br />
<br />
[[File:mem_njit_FW50M_dLvsdD_20-30.jpg|500px|thumb|right|Figure 7:<br />
Changes in position (&Delta;D) and length (&Delta;L measured at 50% of the peak intensity) between 20 and 30 keV for the coronal sources for the 6 flares and 12 time intervals used in the study.]]<br />
<br />
We carried out the same type of analysis of all flares previously analyzed as primarily coronal hard X-ray sources. <br />
This constituted 13 flares with a total of 26 time intervals. <br />
We eliminated some of these events since we could not be sure for various reasons that the available information was consistent with such a dense loop model. <br />
This left just 6 flares and 12 time intervals where we could be reasonably sure that we were getting a reliable measure of the coronal source length as a function of energy. <br />
Values of the change in position (&Delta;D) are plotted in Figure 7 vs. the change in source length (&Delta;L) for all 12 time intervals between 20 and 30 keV. <br />
The change in length ranges between +5 to -8 arcsec but with an average decrease of 1 arcsec and a standard deviation of 3.5 arcsec. <br />
In no case is it as large as the >10 arcsec as previously reported Ref. [2], {3], [4], [5], and [6]). <br />
<br />
The change in position of the peak between 20 and 30 keV ranges between -0.7 and +4.2 arcsec. Given the location of the flares on the solar disc, this is always consistent with either no change or an increase in altitude. Correcting for the foreshortening effect close to the limb, the source altitudes averaged 20 Mm with a scatter deviation of 9 Mm. <br />
Also, we could determine the height of the source above an assumed vertical semicircular loop between the footpoints. This averaged 10 Mm with a scatter deviation of 7 Mm. <br />
If the source emission is thermal, this altitude location and variation with energy is consistent with the standard CHSKP flare model in which magnetic reconnection in a coronal current sheet results in new hot loops being formed at progressively higher altitudes. <br />
The explanation for the nonthermal emission remains unclear and will require further work.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2018arXiv180904631D "Coronal hard X-ray sources revisited"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2004ApJ...603L.117V "A Coronal Thick-Target Interpretation of Two Hard X-Ray Loop Events"]<br />
<br />
[3] [http://iopscience.iop.org/article/10.1086/524184/meta "RHESSI Hard X-Ray Imaging Spectroscopy of Extended Sources and the Physical Properties of Electron Acceleration Regions in Solar Flares"]<br />
<br />
[4] [https://ui.adsabs.harvard.edu//#abs/2012A&A...543A..53G/abstract "Determination of the acceleration region size in a loop-structured solar flare"]<br />
<br />
[5] [https://ui.adsabs.harvard.edu//#abs/2012ApJ...755...32G/abstract "Properties of the Acceleration Regions in Several Loop-structured Solar Flares"]<br />
<br />
[6] [https://ui.adsabs.harvard.edu//#abs/2013ApJ...766...28G/abstract "The Specific Acceleration Rate in Loop-structured Solar Flares — Implications for Electron Acceleration Models"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Data-driven_radiative_hydrodynamic_modeling_of_SOL2014-03-29Data-driven radiative hydrodynamic modeling of SOL2014-03-292018-09-21T14:58:29Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Data-driven radiative hydrodynamic modeling of SOL2014-03-29 (X1)<br />
|number = 274 <br />
|first_author = Fatima Rubio da Costa <br />
|second_author = <br />
|publish_date = 10 May 2016 <br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::275]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::273]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
A [https://en.wikipedia.org/wiki/Solar_flare solar flare] <br />
results from a sudden energy release involving <br />
[https://en.wikipedia.org/wiki/Magnetic_reconnection magnetic reconnection]<br />
in a domain containing free energy. <br />
A large fraction of the radiated energy originates in the chromosphere, so <br />
studying the chromospheric flare emission is therefore a key for understanding<br />
how the flare energy is transported and dissipated.<br />
<br />
In this Nugget we describe the flare of [[Has event date:: March 29, 2014 17:45]], <br />
SOL2014-03-29 (X1), <br />
from the point of view of the excellent chromospheric observations<br />
now becoming available both from ground-based data and new satellite data.<br />
Our ground-based data come from the<br />
[http://nsosp.nso.edu Sacramento Peak Observatory] in New Mexico,<br />
where the Dunn Solar Telescope makes observations with the<br />
[https://www.arcetri.astro.it/science/solare/IBIS_web.html IBIS] <br />
imaging spectrometer, observing in this case <br />
[http://www.astronomyknowhow.com/hydrogen-alpha.htm H-alpha] <br />
and the Ca II 8542 &#8491; line.<br />
These lines,from singly-ionized ions, form at different heights and cover the <br />
whole chromosphere. <br />
They thus provide excellent diagnostics for the flare response of the<br />
lower solar atmosphere. <br />
<br />
== Estimation of the electron distribution ==<br />
<br />
The analysis reported here [1] assumes the <br />
[http://solarphysics.livingreviews.org/open?pubNo=lrsp-2008-1&amp;page=articlesu7.html electron beam]<br />
("thick-target") model, wherein<br />
electrons accelerated in the corona precipitate into the lower<br />
atmosphere, heating it.<br />
We model the temporal evolution of the electron heating <br />
self-consistently from the integrated X-ray spectra obtained by<br />
RHESSI. <br />
The spectra were fitted to a <br />
[https://de.wikipedia.org/wiki/Bremsstrahlung thermal] component plus a<br />
thick-target, non-thermal component. <br />
To estimate the electron flux, we divided the power of non-thermal <br />
electrons by the cross-sectional area of the footpoints in IRIS <br />
images (2796 &#8491;).<br />
Figure 1 shows these parameters.<br />
<br />
[[File:274f1.png|400px|thumb|center|Figure 1: <br />
Temporal evolution of the non-thermal electron spectral parameters<br />
obtained from fitting the RHESSI spectra (a)-(c); the area of the<br />
footpoint emission in IRIS 2796 &#8491; (d) and the non-thermal electron<br />
beam intensity (e). <br />
The dashed line indicates the time of the maximum integrated<br />
X-ray flux from RHESSI.]]<br />
<br />
The RADYN radiative-hydrodynamics code [2,3] has done the hard work of following <br />
how the atmosphere responds to the energy deposited by the non-thermal<br />
electrons. <br />
We constructed a multi-threaded flare loop model and<br />
used the electron flux inferred from RHESSI as the input to the<br />
radiative hydrodynamic code RADYN to simulate the atmospheric<br />
response. <br />
The temporal evolution of the threads was estimated from<br />
the temporal derivative of the GOES 1-8 &#8491; light curve, assuming that<br />
each spike corresponds to a single burst. <br />
<br />
== The chromospheric emissions ==<br />
<br />
RADYN estimates the emission in H-alpha and Ca II 8542 &#8491;; to study other lines, the atmospheric parameters resulting from RADYN go into a radiative-transfer model for other ions, such as our Mg II signatures.<br />
Comparing the model results to the observations, we find that the H-alpha and Ca II profiles fit the observed line shapes, <br />
while the Mg II IRIS profiles are broader in the wings than the synthetic ones. <br />
We find that the H-alpha and Ca II profiles fit the observed line shapes, <br />
while the Mg II IRIS profiles are broader in the wings than the synthetic ones.<br />
Figure 2 shows some of these profiles.<br />
<br />
[[File:274f2.png|700px|thumb|center|Figure 2:<br />
he H-alpha and Ca II 8542 Å line profiles synthesized from RADYN (black solid line) and observed by IBIS at 17:46:13 UT+18 s and 17:45:54 UT+18 s, respectively, <br />
and at different locations along the ribbons of the flare (purple for the southern ribbon and blue for the northern ribbon) and outside (green). <br />
The black dotted line is the IBIS quiet Sun profile calibrated to RADYN.<br />
]]<br />
<br />
== Conclusions ==<br />
<br />
In general, the synthetic intensities agree with the observed ones.<br />
The simulated MgII h&k line profiles have narrower wings than the<br />
observed ones. This discrepancy can be reduced by using a higher<br />
microturbulence velocity (27 km/s) in a narrow chromospheric layer.<br />
An increase of electron density in the upper chromosphere within a<br />
narrow height range of 800 km below the transition region can turn<br />
the simulated MgII line core into emission and thus reproduce the<br />
single peaked profile, which is a common feature seen in all IRIS<br />
flares. <br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2016arXiv160304951R "Data-driven Radiative Hydrodynamic Modeling of the 2014 March 29 X1.0 Solar Flare"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/1995ApJ...440L..29C "Does a nonmagnetic solar chromosphere exist?"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2005ApJ...630..573A "Radiative Hydrodynamic Models of the Optical and Ultraviolet Emission from Solar Flares"]<br />
<br />
[[Has observation by:: Dunn Solar Telescope IBIS| ]] <br />
[[Has observation by:: RHESSI| ]]<br />
[[Has observation by:: IRIS| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Electron_acceleration_and_hard_X-ray_emission_from_SOL2013-11-09Electron acceleration and hard X-ray emission from SOL2013-11-092018-09-21T14:52:07Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Electron Acceleration and Hard X-ray Emission from SOL2013-11-09<br />
|number = 273 <br />
|first_author = Yuri Tsap<br />
|second_author = Galina Motorina<br />
|publish_date = 5 May 2016 <br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::274]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::272]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
The acceleration of high-energy particles has long been recognized as <br />
one of the fundamental characteristics of a solar flare.<br />
<br />
Until now it has not been clear exactly where the electron acceleration region<br />
is located during the flare energy release. <br />
It is usually considered that electrons are accelerated in the <br />
[https://en.wikipedia.org/wiki/Corona solar corona]. <br />
In particular, they can be accelerated in the region of the top<br />
of a [https://en.wikipedia.org/wiki/Coronal_loop magnetic loop] structure<br />
(models with a compact acceleration region) or in the whole loop<br />
(models with an extended acceleration region). <br />
In our view, we can shed light on this issue based on the relationship <br />
between hard X-ray emissions from the coronal part (looptop) of a flare loop<br />
and its footpoints.<br />
<br />
This Nugget is devoted to the analysis of the hard X-ray emission<br />
from different parts of a flare loop based on RHESSI observations<br />
for the well-studied solar flare of [[Has event date::2013-Nov-09 06:26:09]], <br />
[http://sprg.ssl.berkeley.edu/~tohban/browser/?show=grth1+grth3+qlpcr+qli02+qli04+fergo+rms4a&date=20131109&time=062609&bar=1 SOL2013-11-09 (C1.2)] [1,2].<br />
<br />
== Relationships between hard X-ray emissions from different parts of a flare loop ==<br />
<br />
The hard X-ray images (Figure 1) of the solar flare SOL2013-11-09 (C2.6)<br />
show three sources [1]. <br />
One of them was located between two others and characterized by the <br />
strong hard X-ray emission dominated in some time intervals. <br />
This appears to correspond to the looptop.<br />
<br />
[[File:273f1.png|700px|thumb|center|Figure 1: <br />
SDO/AIA 171 &#8491; image (in inverted colors)<br />
overlaid with RHESSI images at 4.1-4.9~keV (white) near the peak<br />
of the impulsive phase, with contours at 50, 60, 70 and 90% of the<br />
maximum of each image, and at 23.0-27.5 keV (blue, levels 50,<br />
60, 70, and 90% of the image maximum) integrated for ~300 s, <br />
covering the entire main impulsive phase (time intervals<br />
indicated in the figure legend).<br />
]]<br />
<br />
The RHESSI imaging appears to resolve this looptop source as a<br />
dense structure [1], [2], with estimates of n<sub>e</sub> ~ 3 x 10<sup>11</sup> cm<sup>-3</sup><br />
and a characteristic size L ~ 9-18 arcsec.<br />
These estimates apply at the time of maximum of the hard X-ray emission (06:25:41~UT) <br />
in terms of the simultaneous <br />
[http://global.jaxa.jp/projects/sat/solar_b/ Hinode]/[http://solarb.mssl.ucl.ac.uk/SolarB/Solar-B.jsp EIS]<br />
observations [1].<br />
The high plasma density suggests that accelerated<br />
electrons must lose the kinetic energy quite rapidly while<br />
propagating from the looptop to the footpoints. <br />
In particular, supposing the<br />
coronal loop column density N = nL/2 = 1-2 x 10<sup>20</sup> cm<sup>-2</sup>, and that<br />
the cosine between the direction of the<br />
magnetic field and the electron velocity &mu; = 0.5, we find that<br />
accelerated electrons will be collisionaly stopped in the coronal<br />
part of a flare loop if their energy E < E<sub>loop</sub> ~ 10 (N<sub>19</sub>/&mu;)<sup>1/2</sup> = 45-63 keV,<br />
where N<sub>19</sub> = N/10<sup>19</sup><br />
<br />
The relationship between the spectral fluxes of hard X-ray<br />
emission at the coronal part I<sub>lp</sub> and the footpoints I<sub>fp</sub>,<br />
as predicted by the<br />
[http://solarphysics.livingreviews.org/open?pubNo=lrsp-2008-1&amp;page=articlesu7.html collisional thick target] model, can be<br />
represented as [3]<br />
<br />
[[File:273f2.png|500px|thumb|center]]<br />
<br />
where B(a, b, c) is an incomplete<br />
[https://en.wikipedia.org/wiki/Beta_function beta function] and &delta; is the <br />
spectral index of the electron flux. <br />
Then, adopting delta = 5.6 and &epsilon;/E<sub>loop</sub> = 0.5, we<br />
find I<sub>lp</sub>/ I<sub>fp</sub> ~ 11.<br />
The estimate thus obtained contradicts the<br />
observed intensities of the hard X-ray sources in the energy range<br />
23.0-27.5 keV (see Fig.1) and demonstrates that the standard<br />
solar flare model requires at least some modifications.<br />
<br />
== Conclusions ==<br />
<br />
The discrepancy between the model estimates and hard X-ray<br />
observations suggests that electron acceleration can occur not<br />
only in the coronal part of a loop but also in the footpoints.<br />
This inference agrees well with the idea of re-acceleration [4]<br />
and in addition with the solar flare model proposed by Zaitsev and<br />
Stepanov [5]; see also [6]. <br />
Note that an extended acceleration region in the corona does not<br />
significantly improve the relationship (1) since the generation of<br />
the hard X-ray emission in the coronal part will be more effective<br />
in this case.<br />
<br />
== References ==<br />
<br />
# [http://adsabs.harvard.edu/abs/2015A%26A...577A..68S "Direct observation of the energy release site in a solar flare by SDO/AIA, Hinode/EIS, and RHESSI"]<br />
# [http://adsabs.harvard.edu/abs/2015SoPh..290.3573S "Impulsive heating of solar flare ribbons above 10 MK"]<br />
# [http://adsabs.harvard.edu/abs/2004ApJ...603L.117V "A coronal thick-target interpretation of two hard X-ray loop events"]<br />
# [http://adsabs.harvard.edu/abs/2009A%26A...508..993B "Local re-acceleration and a modified thick target model of solar flare electrons"]<br />
# [http://adsabs.harvard.edu/abs/2015SoPh..290.3559Z "Particle acceleration and plasma heating in the chromosphere"]<br />
# [http://adsabs.harvard.edu/abs/1998ARep...42..275T "The stochastic acceleration of upper chromospheric electrons"]<br />
<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has observation by:: SDO AIA| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Extreme_events,_stellar_evolution,_and_magnetic_reconnectionExtreme events, stellar evolution, and magnetic reconnection2018-09-21T14:48:08Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Extreme events, stellar evolution, and magnetic reconnection <br />
|number = 272 <br />
|first_author = Hugh Hudson<br />
|second_author = <br />
|publish_date = 30 April 2016 <br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::273]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::271]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
With RHESSI we generally study the most punctuated form of solar<br />
activity, namely the flares, but these occur in swarms. <br />
The swarms of flares include many tiny ones, as well as some powerful ones,<br />
and recently the latter have become especially interesting in terms<br />
of [http://kepler.nasa.gov Kepler] photometry <br />
(1; see [[High Dispersion Spectroscopy of solar-type superflare_stars| earlier Nugget]])<br />
and remarkable results on radioisotope structures in tree rings [2].<br />
In the solar context, the best we can do for "extreme event" is the <br />
celebrated Carrington flare [3], about which we have also had an<br />
[[The 1859 Space Weather Event Revisited| earlier RHESSI Science Nugget]].<br />
<br />
== Extreme events and solar variability ==<br />
<br />
With this term we mean the "Black Swan" events that seem impossible<br />
(such as the appearance of a black swan) until they appear unpredictably<br />
(as when Europeans spotted them at last in Australia). <br />
In our context it means that one of those Kepler superflares could pop off<br />
on our mild Sun, not very frequently, but with devasting consequences if <br />
it did.<br />
<br />
[[File:272f1.png|700px|thumb|center|Figure 1: <br />
Left, what a "superflare" might look like observationally, if only one<br />
had a "supertelescope" (Maehara); right, what the associated spots<br />
might look like (bottom set; Aulanier).<br />
]]<br />
<br />
Figure 1 shows two suprising contributions to our knowledge of solar<br />
extreme events, in work that has been reported only in the last few<br />
years ([1],[2]).<br />
Either of these might imply that the Sun could do something very drastic.<br />
<br />
How do the extreme events fit into the population of flare magnitudes,<br />
which since the 1950s has been known to follow a general power law?<br />
We described one aspect of this, showing off the beauty of the law, in an <br />
[[Chree Analysis for Flares| earlier Nugget]].<br />
One possibility is simply that they extend the power law, such that if<br />
one waited sufficiently one would eventually catch an event like this.<br />
<br />
But the power law is not fixed in amplitude.<br />
With RHESSI we have now had more than one 11-year cycle to study flare<br />
occurrence, as partly described in the Nugget:<br />
[[Cycle 24 - don't panic yet!]]<br />
Instead, the intercept of the power law slides up and down as activity<br />
waxes and wanes.<br />
For the present-day Sun, we see this amplitude, measured logarithmically<br />
as a maximum-to-minimum ratio, to exceed three decades on a time scale<br />
of 243 days (convenient multiple of a mean rotation period).<br />
Figure 2 shows this with the same data shown on different binnings<br />
of 3, 9, and 27 times 27 days.<br />
Even at the longest binning, there are zeroes in the histogram. <br />
At the depths of solar minimum in 2008, especially, it appears as though<br />
the Sun simply turned off.<br />
This mainly shows that a single power law is too simplistic a descriptor<br />
of solar activity.<br />
There is a physical limitation, imposed by observing time, on the upper limit to the <br />
unobservable maximum-to-minimum ratio.<br />
<br />
[[File:272f2.png|800px|thumb|center|Figure 2: <br />
The solar-cycle variation (25 years) of M-class flare occurrence, as captured by <br />
[GOES}.<br />
The display shows both linear and log representations, at different innings, and the suggestion here is that<br />
the maximum-to-minimum ratio may really diverge.<br />
]]<br />
<br />
== Stellar evolution ==<br />
<br />
So, how does solar/stellar activity evolve, based on flare signatures, <br />
on much longer time scales?<br />
That is the question posed by the Kyoto group of K. Shibata, who have <br />
pioneered so much in Kepler "superflare" research [1] and who generated<br />
the left panel of Figure 1.<br />
<br />
The Sun, at its present comfortable evolutionary stage, is winding its way<br />
along towards a relatively well-understood end: gross expansion into a red<br />
giant, some meandering development as its fuel is spent, and then a drastic <br />
splash into a lovely [planetary nebula] within which a white-dwarf core <br />
remains as the only stellar object left.<br />
Figure 3 shows some planetary nebulae, which indeed take on a wide variety<br />
of forms.<br />
This variety does not seem to be well understood, but the basic points<br />
regarding the evolution certainly are.<br />
We're watching this development from the halfway point, about 5 billion<br />
years from the start and equally far away from the end.<br />
Eventually the Sun will become a planetary nebula such as the celebrated<br />
[https://en.wikipedia.org/wiki/Red_Rectangle_Nebula "Red Rectangle"] [4] or something like<br />
what are pictured in Figure 3 here.<br />
<br />
[[File:272nebulae.png|800px|thumb|center|Figure 3: <br />
Planetary nebulae, courtesy Wikipedia.<br />
]]<br />
<br />
<br />
In a recent <br />
[http://www.nasa.gov/image-feature/goddard/2016/hubble-frames-a-unique-red-rectangle ESA/Hubble Space Telecope/NASA] press release, the physical<br />
issues involved with tracking solar activity may have clarified to a<br />
certain extent. <br />
It is reasonably clear from Figure 3 that the terminal state of solar<br />
development, if it resembles the Red Rectangle, can be nothing more or less<br />
than the result of 2D magnetic reconnection.<br />
<br />
[[File:272f3.jpg|400px|thumb|center|Figure 4: <br />
Magnetic reconnection at the end of the Sun's active lifetime?<br />
Photo courtesy ESA/Hubble and NASA showing the star <br />
[https://en.wikipedia.org/wiki/Red_Rectangle_Nebula HD 44179] and its surrounding planetary nebula.<br />
]]<br />
<br />
== Conclusion ==<br />
<br />
Jokes aside, it is interesting for us to try to understand how extreme forms of stellar magnetic activity may develop.<br />
This also has very practical consequences, since our newly developed human technological civilization may well have great<br />
instabilities that could be brought about by a <br />
[https://www.youtube.com/watch?v=FW8Rsil0lqk solar Black Swan event].<br />
Altough in our discussion of the M-class flare statistics we tried to establish that we are reallly ignorant about these things,<br />
based on a minuscule glimpse of how the Sun is working, the <br />
data [1] and the tree-ring data [2] should make us think twice.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2012Natur.485..478M "Superflares on solar-type stars"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2012Natur.486..240M "A signature of cosmic-ray increase in AD 774-775 from tree rings in Japan"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1859MNRAS..20...13C "Description of a Singular Appearance seen in the Sun on September 1, 1859"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/1975ApJ...196..179C "The peculiar object HD 44179 ('The red rectangle')"]<br />
<br />
[[Has article subject:: stellar flares| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Radio_polarization_signatures_in_twisted_flare_loopsRadio polarization signatures in twisted flare loops2018-09-21T13:25:14Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Radio polarization signatures in twisted flare loops <br />
|number = 271 <br />
|first_author = Ivan Sharykin <br />
|second_author = Alexei Kuznetsov<br />
|publish_date = 25 April 2016 <br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::272]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::270]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
The RHESSI view of a solar flare, via hard X-rays and gamma-rays, is quite<br />
direct in one sense, but quite indirect in another: these radiations<br />
can only come from a [http://solarphysics.livingreviews.org/open?pubNo=lrsp-2008-1&amp;page=articlesu7.html dense target medium].<br />
This basic problem means that model understanding of the high-energy<br />
emissions must include an understanding of the particle transport.<br />
Specifically, one cannot "see" the acceleration site if it has low density,<br />
which might be an important requirement for a successful theory of <br />
particle acceleration in a solar plasma atmosphere.<br />
<br />
By contrast, <br />
[http://www.oxfordreference.com/view/10.1093/acref/9780199609055.001.0001/acref-9780199609055-e-1593 gyrosynchrotron] <br />
radio emissions do not require the presence <br />
of a dense target medium. <br />
Instead, the emission process requires the existence of a magnetic field,<br />
of course, but the distribution of magnetic field and plasma have very<br />
different equilibrium processes.<br />
For plasma or gas, hydrostatic equilibrium can dominate; at a minimum the<br />
solar atmosphere typically has a dense chromosphere and a tenuous corona.<br />
Where the magnetic field dominates, on the other hand, the equilibrium<br />
condition would be essentially an isobaric one (a smoothly varying magnetic pressure,<br />
a state that could be termed "isomagnetobaric"), where<br />
only the [https://en.wikipedia.org/wiki/Maxwell_stress_tensor Maxwell stress tensor] should define the large-scale<br />
structure.<br />
Recall that we believe these flaring regions to have low [https://en.wikipedia.org/wiki/Beta_(plasma_physics) plasma beta], i.e. that magnetic pressure dominates gas pressure.<br />
This means that the gyrosynchrotron emissivity, all else being equal,<br />
will have a smoother distribution in space than (say) the hard X-ray<br />
emissivity.<br />
<br />
What does this imply for an observable, such as the polarization pattern,<br />
of a coronal magnetic structure containing fast electrons?<br />
We have thus carried out numerical experiments [1] using the ''GX Simulator''<br />
software package [2], finding that interesting specific signatures<br />
may appear.<br />
New microwave polarimetric imaging facilities such as <br />
[http://www.ovsa.njit.edu E-OVSA] (see also our <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_New_Day_Dawns earlier Nugget]) and <br />
[http://www.skyandtelescope.com/astronomy-news/astro-sightseeing-in-innermongolia/ Mingantu] are now appearing, and our experiments suggest that the <br />
polarization patterns will be informative about not just the energy-loss<br />
domains (hard X-rays) but possibly also the acceleration regions<br />
of high-energy particles.<br />
<br />
== The Simulations ==<br />
<br />
These simulations consist of analytic approximations to the behavior <br />
of a specified particle population in a fixed magnetic geometry<br />
This is far from realistic in the sense that it does not describe the<br />
dynamics of energy exchange between particles and field, which we know to be important<br />
in the flare process.<br />
One could also study the polarization distribution in more complete<br />
models, but the present ones serve to illustrate at least one interesting conclusion (described below).<br />
Our magnetic model follows that designed by Titov and Demoulin [3] in a <br />
commonly-used description of [https://en.wikipedia.org/wiki/Force-free_magnetic_field non-potential magnetic field] <br />
structures in the solar corona.<br />
<br />
[[File:271f1.png|500px|thumb|center|Figure 1: <br />
Titov-Demoulin models of coronal magnetic fields with small twist <br />
(left column) and large twist (right column).<br />
]]<br />
<br />
In these models the fields colored in magenta are the "strapping" fields<br />
introduced to confine the twisted structure via the magnetic tension, one<br />
component of the <br />
[https://en.wikipedia.org/wiki/Maxwell_stress_tensor Maxwell stress tensor].<br />
The gold field domain connects <br />
opposite polarities across the photospheric neutral line (the line defined by the locus of zero vertical field component) and carries<br />
a current to generate the twist.<br />
We imbed the particle fluxess in these structures, specifying<br />
pitch-angle distributions (relative to the direction of the magnetic field), plasma properties such as density and<br />
temperature, and particle sources.<br />
Then we "turn the crank" on ''GX Simulator'' to see what happens.<br />
<br />
[[File:271f2.png|500px|thumb|center|Figure 2:<br />
Results of ''GX Simulator'' runs on the fields described in Figure 1,<br />
with particle inputs as described in [1].<br />
The alternation of polarization sense clearly appears, correlated<br />
with the twist of the flux-rope model.<br />
]]<br />
<br />
The advantage of these numerical experiments is that one can vary<br />
the parameters, which one cannot do with observations - we observe only<br />
what Nature gives us, and sometimes not very well!<br />
Thus one can make possible discoveries in the simulations just as one can make discoveries<br />
in real observations.<br />
These need to confirmed by theoretical and observational work in any case, but the initial germ of discovery may be in one of<br />
these numerical experiments.<br />
<br />
== Conclusions ==<br />
<br />
The main conclusions from our paper [1] reads as <br />
<br />
"• Inversion of the polarization sign of the radio emission, generated by non-<br />
thermal electrons in the twisted loop located in the center of the solar disk,<br />
has form of the line inclined relatively to the loop axis. Polarization of <br />
radio emission from twisted loop on the solar limb experiences change of <br />
its sign along its axis."<br />
<br />
This experiment predict slanting lines of polarity inversion that cut across the flux-rope structure, which may or may not<br />
be recognizable in direct images.<br />
Beyond this greater complexities can be imagined and explored with experiments of this type.<br />
<br />
== References ==<br />
<br />
[1] [http://arxiv.org/abs/1604.05618 "Modelling of Nonthermal Microwave Emission From Twisted Magnetic Loops"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2012AAS...22020451N Integrated IDL Tool For 3d Modeling And Imaging Data Analysis]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1999A%26A...351..707T "Basic topology of twisted magnetic configurations in solar flares"]<br />
<br />
[[Has article subject:: simulations| ]]<br />
[[Has article subject:: radio| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Radio_polarization_signatures_in_twisted_flare_loopsRadio polarization signatures in twisted flare loops2018-09-21T13:24:36Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Radio polarization signatures in twisted flare loops <br />
|number = 271 <br />
|first_author = Ivan Sharykin <br />
|second_author = Alexei Kuznetsov<br />
|publish_date = 25 April 2016 <br />
|next_nugget = Extreme events, stellar evolution, and magnetic reconnection <br />
|previous_nugget = [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/An_Unreported_White-light_Prominence A previously unreported white-light prominence]<br />
}}<br />
<br />
== Introduction ==<br />
<br />
The RHESSI view of a solar flare, via hard X-rays and gamma-rays, is quite<br />
direct in one sense, but quite indirect in another: these radiations<br />
can only come from a [http://solarphysics.livingreviews.org/open?pubNo=lrsp-2008-1&amp;page=articlesu7.html dense target medium].<br />
This basic problem means that model understanding of the high-energy<br />
emissions must include an understanding of the particle transport.<br />
Specifically, one cannot "see" the acceleration site if it has low density,<br />
which might be an important requirement for a successful theory of <br />
particle acceleration in a solar plasma atmosphere.<br />
<br />
By contrast, <br />
[http://www.oxfordreference.com/view/10.1093/acref/9780199609055.001.0001/acref-9780199609055-e-1593 gyrosynchrotron] <br />
radio emissions do not require the presence <br />
of a dense target medium. <br />
Instead, the emission process requires the existence of a magnetic field,<br />
of course, but the distribution of magnetic field and plasma have very<br />
different equilibrium processes.<br />
For plasma or gas, hydrostatic equilibrium can dominate; at a minimum the<br />
solar atmosphere typically has a dense chromosphere and a tenuous corona.<br />
Where the magnetic field dominates, on the other hand, the equilibrium<br />
condition would be essentially an isobaric one (a smoothly varying magnetic pressure,<br />
a state that could be termed "isomagnetobaric"), where<br />
only the [https://en.wikipedia.org/wiki/Maxwell_stress_tensor Maxwell stress tensor] should define the large-scale<br />
structure.<br />
Recall that we believe these flaring regions to have low [https://en.wikipedia.org/wiki/Beta_(plasma_physics) plasma beta], i.e. that magnetic pressure dominates gas pressure.<br />
This means that the gyrosynchrotron emissivity, all else being equal,<br />
will have a smoother distribution in space than (say) the hard X-ray<br />
emissivity.<br />
<br />
What does this imply for an observable, such as the polarization pattern,<br />
of a coronal magnetic structure containing fast electrons?<br />
We have thus carried out numerical experiments [1] using the ''GX Simulator''<br />
software package [2], finding that interesting specific signatures<br />
may appear.<br />
New microwave polarimetric imaging facilities such as <br />
[http://www.ovsa.njit.edu E-OVSA] (see also our <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_New_Day_Dawns earlier Nugget]) and <br />
[http://www.skyandtelescope.com/astronomy-news/astro-sightseeing-in-innermongolia/ Mingantu] are now appearing, and our experiments suggest that the <br />
polarization patterns will be informative about not just the energy-loss<br />
domains (hard X-rays) but possibly also the acceleration regions<br />
of high-energy particles.<br />
<br />
== The Simulations ==<br />
<br />
These simulations consist of analytic approximations to the behavior <br />
of a specified particle population in a fixed magnetic geometry<br />
This is far from realistic in the sense that it does not describe the<br />
dynamics of energy exchange between particles and field, which we know to be important<br />
in the flare process.<br />
One could also study the polarization distribution in more complete<br />
models, but the present ones serve to illustrate at least one interesting conclusion (described below).<br />
Our magnetic model follows that designed by Titov and Demoulin [3] in a <br />
commonly-used description of [https://en.wikipedia.org/wiki/Force-free_magnetic_field non-potential magnetic field] <br />
structures in the solar corona.<br />
<br />
[[File:271f1.png|500px|thumb|center|Figure 1: <br />
Titov-Demoulin models of coronal magnetic fields with small twist <br />
(left column) and large twist (right column).<br />
]]<br />
<br />
In these models the fields colored in magenta are the "strapping" fields<br />
introduced to confine the twisted structure via the magnetic tension, one<br />
component of the <br />
[https://en.wikipedia.org/wiki/Maxwell_stress_tensor Maxwell stress tensor].<br />
The gold field domain connects <br />
opposite polarities across the photospheric neutral line (the line defined by the locus of zero vertical field component) and carries<br />
a current to generate the twist.<br />
We imbed the particle fluxess in these structures, specifying<br />
pitch-angle distributions (relative to the direction of the magnetic field), plasma properties such as density and<br />
temperature, and particle sources.<br />
Then we "turn the crank" on ''GX Simulator'' to see what happens.<br />
<br />
[[File:271f2.png|500px|thumb|center|Figure 2:<br />
Results of ''GX Simulator'' runs on the fields described in Figure 1,<br />
with particle inputs as described in [1].<br />
The alternation of polarization sense clearly appears, correlated<br />
with the twist of the flux-rope model.<br />
]]<br />
<br />
The advantage of these numerical experiments is that one can vary<br />
the parameters, which one cannot do with observations - we observe only<br />
what Nature gives us, and sometimes not very well!<br />
Thus one can make possible discoveries in the simulations just as one can make discoveries<br />
in real observations.<br />
These need to confirmed by theoretical and observational work in any case, but the initial germ of discovery may be in one of<br />
these numerical experiments.<br />
<br />
== Conclusions ==<br />
<br />
The main conclusions from our paper [1] reads as <br />
<br />
"• Inversion of the polarization sign of the radio emission, generated by non-<br />
thermal electrons in the twisted loop located in the center of the solar disk,<br />
has form of the line inclined relatively to the loop axis. Polarization of <br />
radio emission from twisted loop on the solar limb experiences change of <br />
its sign along its axis."<br />
<br />
This experiment predict slanting lines of polarity inversion that cut across the flux-rope structure, which may or may not<br />
be recognizable in direct images.<br />
Beyond this greater complexities can be imagined and explored with experiments of this type.<br />
<br />
== References ==<br />
<br />
[1] [http://arxiv.org/abs/1604.05618 "Modelling of Nonthermal Microwave Emission From Twisted Magnetic Loops"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/2012AAS...22020451N Integrated IDL Tool For 3d Modeling And Imaging Data Analysis]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1999A%26A...351..707T "Basic topology of twisted magnetic configurations in solar flares"]<br />
<br />
[[Has article subject:: simulations| ]]<br />
[[Has article subject:: radio| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/An_Unreported_White-light_ProminenceAn Unreported White-light Prominence2018-09-21T13:22:52Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = An unreported white-light prominence <br />
|number = 270 <br />
|first_author = Matt Penn<br />
|second_author = Hugh Hudson <br />
|publish_date = 28 March 2016 <br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::271]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::269]]}}<br />
}}<br />
<br />
== What is a "white-light prominence? ==<br />
<br />
A white-light flare, as originally discovered in 1859 by <br />
[http://adsabs.harvard.edu/abs/1859MNRAS..20...13C Carrington] <br />
and<br />
[http://adsabs.harvard.edu/abs/1859MNRAS..20...15H Hodgson],<br />
clearly shows us the powerful energy release at the onset of a solar flare.<br />
In many cases this "impulsive phase" produces most of the radiated <br />
energy of the flare, and it also coincides with mass elevation into the <br />
corona as the<br />
[https://en.wikipedia.org/wiki/Neupert_effect Neupert effect],<br />
and further into the heliosphere as a<br />
[https://www.spaceweatherlive.com/en/help/what-is-a-coronal-mass-ejection-cme coronal mass ejection].<br />
<br />
But the true white-light flare emissions come from the lower solar <br />
atmosphere as footpoint sources [1].<br />
In rare cases, one also sees continuum emission from the low corona, and<br />
in such a case we call the phenomenon a "white-light prominence."<br />
Such a phenomenon can be observed visually as a protuberance above the limb,<br />
with a brightness sufficient to be noticeable.<br />
Until recently, there has been almost no literature on this topic, although<br />
there have been anecdotal reports (for example, a paper on solar gamma-ray<br />
observations [2] described an unpublished observations by J. W. Harvey<br />
and collaborators, at the McMath telescope, of SOL1980-06-21.<br />
<br />
A [http://www.cloudynights.com/topic/518199-something-incredibly-rare-white-light-prominence/ recent description]<br />
of the white-light prominence associated with SOL1998-11-22 called such a <br />
occurrence "incredibly rare."<br />
We pirate their graphics here as Figure~1, noting that it was observed by<br />
[http://trace.lmsal.com TRACE] and discovered in that database by Harry Warren.<br />
<br />
[[File:270f1.png|600px|thumb|center|Figure 1: <br />
The TRACE observation of SOL1998-11-22 <br />
([http://www.cloudynights.com/topic/518199-something-incredibly-rare-white-light-prominence/ "incredibly rare."])<br />
]]<br />
One clear description of such an event appeared in the literature in 1991 [3], <br />
and the general interpretation is a fairly straightforward one: intense <br />
ablation of the chromosphere during a flare can create dense loops. <br />
The electron densities of these regions can be as high as 10<sup>12</sup> cm<sup>-3</sup>.<br />
These can drive H-alpha into emission when on the disk, and shine by Thomson<br />
scattering in the visible continuum, when projected against the dark sky.<br />
<br />
== SOL2005-09-07 (X17) ==<br />
<br />
The purpose of this Nugget is to document another example observed only<br />
visually, and not reported in the literature elsewhere so far as we know.<br />
This event was observed by author Penn, Eric Galayda, Aimee Norton, and<br />
Claude Plymate. <br />
None of them could find a camera lying around, amazingly enough, and so only<br />
visual impressions could be captured:<br />
<br />
Eric:<br />
<br />
It was most certainly a limb event. Started as a "am I seeing this?" moment <br />
but quickly grew, bulging off-limb as it brightened. We were between setups <br />
and really were kicking our bad luck that there was nothing to use to image <br />
the flare. I believe even our cheap digital camera was without batteries. <br />
<br />
Claude:<br />
<br />
I was setting up the NAC for Matt and was in the early process of optical<br />
alignment. To check focus, I climbed up on top of the spectrograph guider<br />
platform. While trying to determine best focus, I noticed and odd "glint"<br />
projecting maybe a couple cm above the side of the solar disk imaged onto<br />
the table. It looked like some stray light reflecting from somewhere. I<br />
used a piece of paper to try and trace where the stray light was coming<br />
from. To my surprise, it seemed to be on axis and part of the image. I<br />
then noticed that the "glint" seemed to have risen a bit further above the<br />
limb. Perplexed, I thought about it for a moment, then yelled over at Eric<br />
to check the GOES X-ray monitor. The plot he pulled up showed an above<br />
X-class flare in progress!!! Holy $#&%!!! That's when I realized that we<br />
were actually seeing a corona mass ejection in white light!<br />
<br />
...an above X-class flare in progress!!! Holy $#&%!!!<br />
<br />
Looking closely at the image, I could see a brightening at the foot point<br />
and was able to discern that the color wasn't quite white but more a pale<br />
purple. My guess at the time was that the purple was coming from a mix of<br />
both H-alpha and H-beta. About then was when Aimee walked in for her tour<br />
of the telescope. As you might imagine, we were all going ape and I can<br />
only guess her first impression of us. By that time, the CME had continued<br />
to move higher, had disconnected from the solar disk and was fading. I<br />
think the entire incident only lasted about 20 minutes.<br />
<br />
Aimee:<br />
<br />
I mostly remember Claude being very animated, literally jumping about <br />
while a dynamic feature writhed just off the limb of the Sun ... I think <br />
Claude is wrong that I walked in on while it was going on. I was there <br />
from the start. I distinctly remember him going through the phase of <br />
thinking it was scattered light, placing the paper down and then starting <br />
to swear. <br />
<br />
Matt:<br />
<br />
I remember [the flare] appearing blue-ish white to me, like a fluorescent <br />
bulb; two bright regions maybe 30-50% brighter than the quiet Sun, maybe <br />
each about 5mm across in size [at 2.4 arc sec/mm scale]. I remember <br />
asking if we had any sort of digital camera around and we didn't... it was in <br />
the days before we carried cell phones with cameras up there I guess. <br />
<br />
We left the image and continued with the tour, but a few minutes later Claude <br />
called us back, saying that there was a prominence visible in white light. <br />
We went back to the projected image and there was a prominence off the limb <br />
of the Sun as plain as day. It was bright, maybe 5-8 cm off the limb and <br />
5-8cm long. It was very thin, but showed structures, in particular, very <br />
bright knots. The knots seemed like point sources, and either were flickering <br />
in the seeing or showed real time evolution. I think the clouds were getting <br />
worse (they were thick cumulus blowing by the summit) and we watched the <br />
prominence for a few more minutes. It stayed visible for that time but was <br />
eventually covered by clouds.<br />
<br />
There are definitely some scientific results in these commentaries, for example<br />
the idea that there might be bright points ("maybe 30-50% brighter than the quiet Sun")<br />
that might not have been resolved in other observations.<br />
Don't forget the large (1.6 m) aperture of the<br />
[https://en.wikipedia.org/wiki/McMath–Pierce_solar_telescope McMath-Pierce]<br />
telescope, with a diffraction limit near 0.1 arc sec.<br />
Also the color remarks ("pale purple" and "blue-ish white, like a fluorescent bulb") <br />
call to mind the comparison of Carrington's flare with &alpha; Lyrae<br />
([https://en.wikipedia.org/wiki/Vega Vega]).<br />
<br />
As regards other data on this fine X17 event of [[Has event date:: September 05, 2005 17:43]]<br />
, Figure 2 shows RHESSI and <br />
[http://lasco-www.nrl.navy.mil LASCO]<br />
file information here.<br />
<br />
[[File:270f2.png|600px|thumb|center|Figure 1: <br />
File images of the event: left, a RHESSI 25-keV image; right, the CME <br />
aftermath (a bright ray and fainter structure_ recorded by the <br />
[https://en.wikipedia.org/wiki/Mauna_Loa Mauna Loa]<br />
[https://www2.hao.ucar.edu/news/2013-aug/end-era-mauna-loa-k-coronometer-decommissioned K-coronameter].<br />
]]<br />
<br />
== Conclusion ==<br />
<br />
This event adds to our historical knowledge, albeit informal, of <br />
white-light prominences.<br />
As the Hiei event showed ([3]), though, and Ref. [4] confirmed, one can<br />
see these things if one looks for them (for example, in the<br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Flare_Coronal_Rain data] from<br />
[http://hmi.stanford.edu HMI]).<br />
They are not in fact "incredibly rare", though they do require exceptional<br />
circumstances.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/1859MNRAS..20...13C "Description of a Singular Appearance seen in the Sun on September 1, 1859"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/1990ApJS...73..213C "Emission characteristics of three intense solar flares observed in cycle 21"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1992PASJ...44...55H "White-light flare observed at the solar limb"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/2014ApJ...780L..28M "Chromospheric and Coronal Observations of Solar Flares with the Helioseismic and Magnetic Imager"]<br />
<br />
[[Has observation by:: LASCO| ]] <br />
[[Has observation by:: RHESSI| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI%27s_5th_AnnealRHESSI's 5th Anneal2018-09-18T15:51:18Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI's 5th Anneal <br />
|number = 269<br />
|first_author = Albert Shih<br />
|second_author = Brian Dennis<br />
|publish_date = 23 February 2016 <br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::270]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::268]]}}<br />
}}<br />
<br />
== The anneal ==<br />
<br />
RHESSI's detectors accumulate radiation damage, and require periodic re-annealing.<br />
This involves bringing the temperatures up to a certain point, whereupon the crystal defects that have built up will tend to heal themselves.<br />
Then a careful cool-down to the normal operating temperature restores the pre-anneal performance.<br />
At least that is the hope, but it is a delicate operation and surprises have happened in past episodes. <br />
We have previously annealed in 2007, 2010, 2012, and 2014, and maintain a table [http://hesperia.gsfc.nasa.gov/ssw/hessi/dbase/rhessi_nosun_times.txt here] that lists not only the anneal time intervals, but also the times of offpoint from the Sun.<br />
These are all of the major data gaps in our solar coverage.<br />
The present one started on February 23, and will last probably until mid-April.<br />
Figure 1 shows the last flare imaged in the current series (see this [http://sprg.ssl.berkeley.edu/~tohban/browser/?show=grth1+qlpcr+qli03+fergo&date=20160222&time=205359&bar=1 Browser] summary): a B8.7 flare that took place on [[Has event date:: Feb 2, 2016 20:55]], SOL2016-02-22T20:55.<br />
<br />
[[File:269f1.png|500px|thumb|center|Figure 1: <br />
The final flare in the current sequence, as shown in the RHESSI flare catalog, courtesy [http://sprg.ssl.berkeley.edu/~tohban/browser/Browser Browser].]]<br />
<br />
For some of the history here, please see these Nuggets:<br />
The [[Annealing RHESSI for the first time]], 2007, some six years after launch.<br />
The [[RHESSI is Annealing Now| fourth anneal]], 2014; and its [[RHESSI has resumed operations| conclusion]].<br />
We will issue another Nugget when operations resume.<br />
<br />
== Annealing... ==<br />
<br />
Figure 2 shows the time history of the annealing procedure, from which a slow return to operating temperature is under way... to be continued....<br />
<br />
[[File:269new3.png|500px|thumb|center|Figure 2: <br />
Time history of the anneal, reflected in the temperature telemetered from the cold plate.]]<br />
<br />
The Sun has cooperated thus far; RHESSI has not missed any spectacular gamma-ray flares (yet).<br />
Figure 3 shows this cooperative spirit in terms of the GOES soft X-ray time history from the beginning of 2016.<br />
The most recent M-class flare was on 15 February.<br />
<br />
[[File:269new2.png|500px|thumb|center|Figure 3: <br />
GOES time history of the anneal, with the beginning of the anneal marked with a dashed line.<br />
Almost no flares at all!]]<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has article subject:: detector anneal| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI%27s_5th_AnnealRHESSI's 5th Anneal2018-09-18T15:40:26Z<p>Schriste: /* The anneal */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI's 5th Anneal <br />
|number = 269<br />
|first_author = Albert Shih<br />
|second_author = Brian Dennis<br />
|publish_date = 23 February 2016 <br />
|next_nugget = A meritorious index <br />
|previous_nugget = [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_new_development_in_the_Frost-Dennis_paradigm The Frost-Dennis paradigm]<br />
}}<br />
<br />
== The anneal ==<br />
<br />
RHESSI's detectors accumulate radiation damage, and require periodic re-annealing.<br />
This involves bringing the temperatures up to a certain point, whereupon the crystal defects that have built up will tend to heal themselves.<br />
Then a careful cool-down to the normal operating temperature restores the pre-anneal performance.<br />
At least that is the hope, but it is a delicate operation and surprises have happened in past episodes. <br />
We have previously annealed in 2007, 2010, 2012, and 2014, and maintain a table [http://hesperia.gsfc.nasa.gov/ssw/hessi/dbase/rhessi_nosun_times.txt here] that lists not only the anneal time intervals, but also the times of offpoint from the Sun.<br />
These are all of the major data gaps in our solar coverage.<br />
The present one started on February 23, and will last probably until mid-April.<br />
Figure 1 shows the last flare imaged in the current series (see this [http://sprg.ssl.berkeley.edu/~tohban/browser/?show=grth1+qlpcr+qli03+fergo&date=20160222&time=205359&bar=1 Browser] summary): a B8.7 flare that took place on [[Has event date:: Feb 2, 2016 20:55]], SOL2016-02-22T20:55.<br />
<br />
[[File:269f1.png|500px|thumb|center|Figure 1: <br />
The final flare in the current sequence, as shown in the RHESSI flare catalog, courtesy [http://sprg.ssl.berkeley.edu/~tohban/browser/Browser Browser].]]<br />
<br />
For some of the history here, please see these Nuggets:<br />
The [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 first anneal], 2007, some six years after launch.<br />
The [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_is_Annealing_Now fourth anneal], 2014; and its <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Resumes_Observations conclusion].<br />
We will issue another Nugget when operations resume.<br />
<br />
== Annealing... ==<br />
<br />
Figure 2 shows the time history of the annealing procedure, from which a slow return to operating temperature is under way... to be continued....<br />
<br />
[[File:269new3.png|500px|thumb|center|Figure 2: <br />
Time history of the anneal, reflected in the temperature telemetered from the cold plate.]]<br />
<br />
The Sun has cooperated thus far; RHESSI has not missed any spectacular gamma-ray flares (yet).<br />
Figure 3 shows this cooperative spirit in terms of the GOES soft X-ray time history from the beginning of 2016.<br />
The most recent M-class flare was on 15 February.<br />
<br />
[[File:269new2.png|500px|thumb|center|Figure 3: <br />
GOES time history of the anneal, with the beginning of the anneal marked with a dashed line.<br />
Almost no flares at all!]]<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has article subject:: detector anneal| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_Sunspot_from_Cycle_25_for_sureA Sunspot from Cycle 25 for sure2018-09-18T11:32:40Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = A Sunspot from Cycle 25 for sure <br />
|number = 321<br />
|first_author = Tomek Mrozek <br />
|second_author = Hugh Hudson<br />
|publish_date = 10 April 2018<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::322]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::320]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
This brief Nugget simply announces that YES, we really have seen <br />
[https://en.wikipedia.org/wiki/Solar_cycle_25 Cycle 25]<br />
[http://www.scholarpedia.org/article/Solar_activity sunspot activity].<br />
An [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_Curious_Sunspot_Group_in_2018 earlier Nugget] <br />
hinted at this, but it was not so clear a case as presented by today's<br />
new tiny sunspot.<br />
Why is this interesting? <br />
It's because spots appearing this early in a cycle - even before a minimum is <br />
well established - are quite rare.<br />
We could speculate that solar minimum may arrive early and/or may be brief,<br />
more evidence regarding the seemingly <br />
[https://en.wikipedia.org/wiki/Stochastic stochastic] component of the <br />
development of the<br />
[https://solarscience.msfc.nasa.gov/SunspotCycle.shtml solar magnetic cycle].<br />
<br />
== The Cycle 25 Sunspot ==<br />
<br />
At the time of writing, <br />
[https://www.swpc.noaa.gov NOAA] has not identified<br />
this new sunspot with an official active-region number, and so <br />
there could be some things to quibble about.<br />
But the <br />
[https://solarscience.msfc.nasa.gov/SunspotCycle.shtml magnetic polarity]<br />
of the region unmistakeably identifies it as a piece of the new cycle,<br />
because it reverses the polarity expected for<br />
[https://en.wikipedia.org/wiki/Solar_cycle_24 Cycle 24] regions.<br />
<br />
Figure 1 here shows the new spot as of this date (10-April-2018).<br />
It is marginally detectable but definitely there in relatively crude<br />
1024x1025 .gif versions of the beautiful data from the <br />
[https://sdo.gsfc.nasa.gov/data/dataaccess.php SDO] space observatory.<br />
<br />
[[File:321f1.png|650px|thumb|center|Figure 1: <br />
File images from the<br />
[http://hmi.lmsal.com HMI] instrument on SDO:<br />
left, the continuum intensity; right, the telltale magnetic field.<br />
From the latter one can see black polarity to the right ("preceding",<br />
as the Sun rotates).<br />
This is the opposite of that shown, for example, by the exceedingly <br />
tiny region at about -5 degrees.<br />
]]<br />
<br />
It requires a bit of patience to see the spot; refer to the location <br />
of the magnetic features and perhaps dither the window on your browser screen.<br />
The icon for this Nugget on the parent page here has a slightly better <br />
view derived from a 4096x4096 image.<br />
<br />
== And a flare! ==<br />
<br />
Sam Freeland notes that this little region arguably has produced the<br />
[http://sdowww.lmsal.com/sdomedia/ssw/media/ssw/ssw_client/data/ssw_service_180409_152748_17806_2/www/ first flare of Cycle 25].<br />
We might label this SOL2018-04-09T12:57 (A2.5).<br />
We thank Bill Marquette and Dick Canfield for relaying this information.<br />
Could such a tiny event have been detected by GOES at the corresponding phase of the previous cycle?<br />
<br />
== Conclusion ==<br />
<br />
This sunspot has been tabulated in the excellent <br />
[http://www.solen.info/solar/cycle25_spots.html SOLEN] page of<br />
Jan Alvestad.<br />
The Nugget-writers here thank him for his thorough monitoring of solar<br />
activity, and also thank Leif Svalgaard for paying close attention as well.<br />
<br />
[[Has observation by:: SDO HMI| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/NuSTAR_detects_X-ray_flares_in_the_quiet_SunNuSTAR detects X-ray flares in the quiet Sun2018-09-18T11:32:01Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = NuSTAR detects X-ray flares in the quiet Sun <br />
|number = 319<br />
|first_author = Matej Kuhar <br />
|second_author = <br />
|publish_date = 26 March 2018<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::320]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::318]]}}}}<br />
<br />
== Introduction ==<br />
<br />
It is still unexplained why the <br />
[https://solarscience.msfc.nasa.gov/corona.shtml solar corona]<br />
(about 1 MK) is so much hotter than the <br />
[https://solarscience.msfc.nasa.gov/surface.shtml photosphere] (about 5700 K). <br />
While it is clear that this energy comes from the Sun's magnetic field<br />
somehow, rival conversion mechanism(s) are still much discussed in the <br />
community. <br />
The <br />
[https://en.wikipedia.org/wiki/Solar_flare solar flare] mechanism <br />
is one of the most promising candidates, since <br />
a flare converts large amounts of magnetic energy into other forms. <br />
Most energetic flares are, however, too rare to explain the steady supply <br />
of energy to the corona, which maintains its high temperature even during <br />
solar minima, when no extreme solar activity takes place. <br />
It has been debated that the flares on the smallest scale might supply <br />
enough energy not through their (relatively small) individual energy <br />
budgets, but through their large numbers. <br />
Numerous studies have shown that solar flares follow a power-law <br />
distribution as a function of energy, with the power-law index <br />
fluctuating around 2. <br />
It is this value that divides the two regimes (Ref. [1]): for index &lt;2, <br />
the largest flares contribute more to coronal heating; for index &gt;2, <br />
the smallest flares are the main contributors. <br />
Parker (Ref. [2]) had proposed the scenario where <br />
[https://en.wikipedia.org/wiki/Nanoflares nanoflares] <br />
(with energies of 10<sup>24</sup> erg or less) make an ensemble of <br />
unresolved basic reconnection events that occupy the whole solar disk <br />
and thus make up the X-ray corona. <br />
<br />
== NuSTAR sees quiet-Sun flares ==<br />
<br />
NuSTAR is a focusing-optics <br />
[https://imagine.gsfc.nasa.gov/science/toolbox/xray_telescopes1.html hard X-ray telescope] operating in the energy range 3-79 keV and <br />
launched by NASA in 2012. <br />
Even though it was not designed for solar observations, it can be <br />
pointed to the Sun. <br />
This capability has led to 12 brief solar observing sessions in the <br />
last 3.5 years, which have resulted in observations of <br />
[https://apod.nasa.gov/apod/ap020801.html active-region] <br />
"microflares," quiescent active regions, high altitude coronal emissions, etc.;<br />
a microflare is just a weaker version of a true solar flare.<br />
<br />
In the July 2016 and March 2017 observing campaigns NuSTAR also observed <br />
three tiny flare-like brightenings in the quiet Sun (QS).<br />
These observations, to the authors' best knowledge, represent <br />
the first spectroscopically resolved X-ray measurements of QS flare events<br />
in hard X-rays (HXR). <br />
They were associated with the solar magnetic <br />
[http://www.scholarpedia.org/article/Solar_activity network] <br />
and, even though feeble when compared to larger <br />
[https://en.wikipedia.org/wiki/Solar_flare M- and X-class] flares, <br />
they showed complex morphologies.<br />
<br />
[[File:319f1.png|500px|thumb|center|Figure 1: <br />
Three quiet-Sun flares observed with NuSTAR. <br />
Left: Flare morphologies. <br />
Middle: Time evolution of NuSTAR and <br />
[http://aia.lmsal.com AIA] EUV fluxes. <br />
Right: Spectra of the observed events. <br />
The green points represent emission during the flare peak, <br />
while pink represents the background. <br />
]]<br />
<br />
== Spectra ==<br />
<br />
Flare-integrated spectra were made for each event, using a spectral model<br />
consisting of an isothermal component plus a fixed background.<br />
The fits provided temperatures in the range 3.1-4.2 MK and emission measures <br />
in the range (0.6-15) x 10<sup>44</sup> cm<sup>-3</sup>. <br />
These values give thermal energy contents in the range <br />
(1.8-6.0) x 10<sup>26</sup> erg and place the observed events (brown box) <br />
just between the usual microflares observed in active regions and smallest <br />
QS flares observed previously in soft X-raus and the EUV <br />
(see the left panel of Figure 2; the brown box shows the NuSTAR events). <br />
The same story is shown in the right panel of the same figure in the<br />
parameter space of temperature and emission measure, where the RHESSI <br />
observation range is shown in yellow, NuSTAR AR microflares in blue, <br />
NuSTAR QS flares in orange, while the faintest events in the soft X-rays <br />
and the EUV observed previously (e.g. Ref. [3]) are shown in green. <br />
The same classes of events in the two plots are highlighted with arrows. <br />
In the [https://www.swpc.noaa.gov/products/goes-x-ray-flux GOES] <br />
classification, our events occupy the range between <br />
1/1000 and 1/100 of a GOES A-class event, the weakest classification. <br />
The spectra did not show direct signs of a nonthermal component. <br />
However, since upper limits of energy in a hidden nonthermal component <br />
(still consistent with the observations, i.e. <br />
no emission above 5 keV) estimated at ~5 x 10<sup>26</sup> erg are within <br />
the uncertainties of thermal energy estimates, and the non-detection of <br />
the nonthermal signal is not a constraining result.<br />
<br />
[[File:319f2.png|600px|thumb|center|Figure 2:<br />
NuSTAR QS events in the context of previous studies. <br />
Left: Frequency distribution of flares as estimated from various <br />
soft X-ray, hard X-ray, and EUV observations. <br />
Right: Different types of flares in temperature-emission measure parameter <br />
space. <br />
The GOES classes are shown with isocurves.<br />
]]<br />
<br />
== Future runs ==<br />
<br />
A fuller account of this work is in Ref. [4].<br />
With the decreasing solar activity towards solar minimum in 2019/2020, <br />
we expect progressively better conditions for NuSTAR QS observations. <br />
We expect higher livetimes and even lower backgrounds in the future runs, <br />
which should allow for observations of even smaller events; we estimate<br />
improvement of a factor of about 70 in sensitivity.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/1991SoPh..133..357H "Solar flares, microflares, nanoflares, and coronal heating"]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/1988ApJ...330..474P "Nanoflares and the solar X-ray corona"]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/1997ApJ...488..499K "X-Ray Network Flares of the Quiet Sun"]<br />
<br />
[4] [http://adsabs.harvard.edu/abs/2018arXiv180308365K "NuSTAR Detection of X-Ray Heating Events in the Quiet Sun"]<br />
<br />
[[Has observation by:: NuSTAR| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_myriad_of_microflaresA myriad of microflares2018-09-18T11:30:30Z<p>Schriste: /* References */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=A myriad of microflares<br />
|number=52<br />
|first_author=Iain Hannah<br />
|second_author=Steven Christe<br />
|publish_date=2 January 2007<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::53]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::51]]}}<br />
}}<br />
<br />
==Introduction==<br />
<br />
RHESSI has been operating for over 5 years now and has seen quite a number of flares: the official flare list (thanks to Jim McTiernan) has nearly 11,000 flare from March 2002 to March 2007. Of these, 7,839 are microflares, that is low C- Class to sub A-Class flares, with the name originally denoting events with energies "micro" or a millionth, of a larger flare. RHESSI is especially good at observing these wee events due to its automated shutters. During large flares these shutters are placed in front of the detectors protecting them from the large fluxes of low energy thermal photons but during quieter times they can be moved out, allowing the full sensitivity of RHESSI's detectors to be revealed.<br />
<br />
The official flare list however looks for events in 12-25 keV, an energy range in which many microflares rarely get above background. So we developed a new algorithm that is optimised to search for microflares in 6-12 keV during shutter out times. We find 24,799 events, over 3 times the number from the official flare list.<br />
<br />
With so many events, there are many nuggets to mine but in order to present a concise nugget we will concentrate on when and where these RHESSI microflares occur.<br />
<br />
==How often?==<br />
<br />
The first question to answer is how often do these events occur.<br />
<br />
[[File:52f1.png|600px|thumb|center|Figure 1: The microflaring rate over the past 5 years of observations. The histogram has been corrected for the times RHESSI was not observing or had the shutter in, hence the rate per live day. The red line and righthand axis indicate the number of sunspots (via NOAA). We can clearly see that the microflaring rate is closely tied to solar activity, with the number of microflares dropping as we approach solar minimum in early 2007.]]<br />
<br />
<br />
In Figure 1, we see that the histogram of the microflaring rate steadily decreases from solar maximum to minimum: in 2002 we could expect over 70 microflares per day, but by solar minimum in 2007 be get under 10 a day on average. Also shown in Figure 1 is the time profile of sunspot number over the same time period. Clearly the microflaring rate is tied to solar activity, which is unsurprising. Although over the same period the same is not true for the largest flares (see a previous nugget). Note the discrepancy between the microflaring rate and the sunspot number during the first 6 months is, we believe, due to a combination of using a different strategy as to when the shutters came in and the large number of C, M and X-class flares hiding smaller events.<br />
<br />
==Where on the solar disk?==<br />
<br />
The next question is to ask where on the solar disk these events occur.<br />
<br />
[[File:52f2.png|600px|thumb|center|Figure 2: The longitude and latitude positions of all the microflares with trusted positions, over 24,000 events.]]<br />
<br />
Flare positions (as seen in Figure 2) reinforce the sense that these microflares are associated with active regions. Their latitudes are clustered in the active region band (where active regions usually appear) and in the longitude plot we can see the distinct straight lines where active regions are traced out by their microflares as the move across the disk. A small number events have positions not associated with active regions but on closer inspection we do not trust the position information for these events, as the RHESSI's roll solution is not correct. So all the 6-12 keV events RHESSI observes are active region phenomena.<br />
<br />
==Imaging using visibilities==<br />
<br />
Of course with RHESSI we can do considerably more than just the times and positions of these events. We can for instance, investigate the image of the thermal emission, which we take to be 4-8 keV in these events. The resulting loop like structures are important as they give us an estimate of the emitting thermal volume and hence thermal energy.<br />
<br />
[[File:52f3.png|600px|thumb|center|Figure 3: Images, using MEM_NJIT on the 4-8 keV visibilities, for the most prolific microflaring active region RHESSI observed, AR10536. The colour scale changes from image to image, so those with the brighter backgrounds are smaller events. The solid circle indicates the size of the sunspot group and the time profile below the image shows the GOES 1-8 &Acirc; light curve for time +/- 12 hours of the microflare. Note that not all the flares as they are either from other active regions or occurred during periods when RHESSI was not observing with shutters out.]]<br />
<br />
In order to quickly investigate the spatial information about so many event we use visibilities (as detailed in a previous nugget). Visibilities allow the spatial structure to be quickly recovered as they make use of the data in the form of a compact set of calibrated data, instead of the full RHESSI un-calibrated time profile. Shown in Figure 3 are the resulting 4-8 keV images using the MEM_NJIT algorithm ( Ref. [1]), for 16 seconds about the peak time in 6-12 keV for each microflare we associated with NOAA active region 10536. We see that some microflares are repeat flaring events from the same loop structure. Others are occurring at different locations throughout the active region, either compact points or large loops reaching out from the active region, with one leg in the sunspot group.<br />
<br />
==A wealth of information==<br />
<br />
In this nugget we have shown that RHESSI microflares are closely related to solar activity and are associated with active regions. But this is just one brief aspect of these events.<br />
<br />
We have a wealth of information about these events with <br />
[https://hesperia.gsfc.nasa.gov/rhessi3/software/imaging-software/forward-fit/index.html forward fitting] shapes to the visibilities allowing us to investigate the spatial scales of the emission and fitting the spectrum of these events allows us to obtain the temperature, emission measure as well as characteristics of the non-thermal component. Combining all this information allows us to investigate the distributions of both the thermal and non-thermal energies in these events and the implications this has for small flare and coronal heating. Something that will be the topic of a future Nugget......<br />
<br />
== Biographical Note ==<br />
Iain Hannah and Steven Christe are both members of the RHESSI team at Space Science Lab, UC Berkeley.<br />
<br />
== References ==<br />
[1] [http://adsabs.harvard.edu/abs/2007SoPh..240..241S "Analysis of RHESSI Flares Using a Radio Astronomical Technique"]<br />
<br />
[[Has article subject:: microflares| ]]<br />
[[Has observation by:: RHESSI| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Annealing_RHESSI_for_the_first_timeAnnealing RHESSI for the first time2018-09-18T11:29:17Z<p>Schriste: /* Biographical note */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Annealing RHESSI for the first time<br />
|number=69<br />
|first_author=David Smith<br />
|publish_date=12 January 2008<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::70]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::68]]}}<br />
}}<br />
<br />
==Introduction==<br />
<br />
The RHESSI detectors are high-purity germanium semiconductor diodes, operated in a unique segmented manner (please see our <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/X-ray_and_gamma-ray_spectroscopy earlier Nugget] for more detailed information). All such detectors are sensitive to radiation damage by high-energy particles in space. In RHESSI's low-Earth orbit, the particles are those of the Van Allen belts, and the dosage can be readily predicted. Thus when RHESSI was launched we were aware that the detectors would have a finite useful lifetime. Gradually, as damage builds up, the detector resolution and effective volume drop.<br />
<br />
Annealing - accomplished by heating the detector to a relatively high temperature (say 100 C) and allowing it to soak at this elevated temperature (say for a week) can restore functionality. The RHESSI team had always planned this but became more confident as the other Ge detector array, that on the gamma-ray observatory [http://sci.esa.int/integral/ Integral], benefited from routine annealing operations.<br />
<br />
[[File:69f1.png|400px|thumb|center|<br />
Figure 1: Cross-section of one of RHESSI's cylindrical Ge detectors, showing the general structure of the electric fields in the detector volume. <br />
]]<br />
<br />
Figure 1 (above) shows the geometry of the RHESSI detectors, along with the pattern of the electric field within the Ge volume that sweeps out the charge left by a photon interaction. <br />
The Sun would be at the top, and the front segment is that part of the volume above the dashed line, the rear segment below. Separate anodes (contacts made in the hollow central region) separately collect the charge from photon interactions the two segments. <br />
Radiation damage to the crystal traps moving charges that are part of the signal from a gamma-ray detection, resulting in a partial loss of signal and poorer energy resolution. It also creates permanent charges throughout the detector volume, which can distort the imposed electric field and, when the damage is particularly severe, create dead volumes within the crystal.<br />
<br />
<br />
==The 2008 RHESSI anneal==<br />
<br />
After much preparation, the RHESSI team annealed the detectors in November for the first time, almost six years after launch. By this time the radiation damage had built up to such a degree that normal operation, especially for gamma-rays, was rapidly becoming difficult. The operation was worrisome, since the segmentation (see Figure 1) might have been destroyed by the anneal. The Integral detectors are not segmented and do not have this extra risk factor. That was the reason to wait so long for the first anneal, and to keep the temperature up only for a limited period of time (one week at 90 C). The result is shown in the figure below - a success roughly as expected, but not without surprises. The anneal did not restore the full resolution of the detectors, but that is of minimal importance. The main thing is that the sensitivity (volume and efficiency) have returned.<br />
<br />
[[File:69f2.png|600px|thumb|center|<br />
Figure 2: Response of the RHESSI rear segments at the 511-keV line of positron annihilation (a background feature). Green was the situation before the anneal, and red the situation afterwards. The black (initial) and blue (mid-2006) behavior shows that the anneal fully restored the detectors' effective volume if not the original energy resolution.<br />
The annealing also restored the front segments to good performance, as shown in Figure 3. Seven individual detectors are shown as separate curves for three solar flares observed in 2005 (left), then just before the anneal (middle), and at present after the anneal (right). Note the disorder and disagreement among the detectors just before (the middle panel) and the good agreement afterwards.<br />
]]<br />
<br />
[[File:69f3.png|600px|thumb|center|<br />
Figure 3: Spectra of three solar flares observed during normal operation (left), just before the anneal (center), and after the anneal (right). These spectra cover the soft X-ray range with a peak at about 7 keV. The radiation damage destroyed the agreement among the 9 detectors (central panel) making analysis very difficult, but the anneal has restored the agreement satisfactorily.<br />
]]<br />
<br />
==The future==<br />
<br />
With this successful annealing operation, RHESSI is ready for the energetic flares expected in Cycle 24, which is just beginning (see our earlier Nugget, in particular its Figure 3). We can do it again if need be, and in fact expect to do so as radiation damage again gradually builds up.<br />
<br />
==Biographical note==<br />
David Smith is a RHESSI team member at UC Santa Cruz.<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has article subject:: detector anneal| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI%27s_5th_AnnealRHESSI's 5th Anneal2018-09-18T11:28:19Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI's 5th Anneal <br />
|number = 269<br />
|first_author = Albert Shih<br />
|second_author = Brian Dennis<br />
|publish_date = 23 February 2016 <br />
|next_nugget = A meritorious index <br />
|previous_nugget = [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/A_new_development_in_the_Frost-Dennis_paradigm The Frost-Dennis paradigm]<br />
}}<br />
<br />
== The anneal ==<br />
<br />
RHESSI's detectors accumulate radiation damage, and require periodic re-annealing.<br />
This involves bringing the temperatures up to a certain point, whereupon the crystal defects that have built up will tend to heal themselves.<br />
Then a careful cool-down to the normal operating temperature restores the pre-anneal performance.<br />
At least that is the hope, but it is a delicate operation and surprises have happened in past episodes. <br />
We have previously annealed in 2007, 2010, 2012, and 2014, and maintain a table [http://hesperia.gsfc.nasa.gov/ssw/hessi/dbase/rhessi_nosun_times.txt here] that lists not only the anneal time intervals, but also the times of offpoint from the Sun.<br />
These are all of the major data gaps in our solar coverage.<br />
The present one started on February 23, and will last probably until mid-April.<br />
Figure 1 shows the last flare imaged in the current series (see this [http://sprg.ssl.berkeley.edu/~tohban/browser/?show=grth1+qlpcr+qli03+fergo&date=20160222&time=205359&bar=1 Browser] summary): a B8.7 flare that took place on [[Has event date:: Feb 2, 2012 20:355]], SOL2012-02-22T20:55.<br />
<br />
[[File:269f1.png|500px|thumb|center|Figure 1: <br />
The final flare in the current sequence, as shown in the RHESSI flare catalog, courtesy [http://sprg.ssl.berkeley.edu/~tohban/browser/Browser Browser].]]<br />
<br />
For some of the history here, please see these Nuggets:<br />
The [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 first anneal], 2007, some six years after launch.<br />
The [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_is_Annealing_Now fourth anneal], 2014; and its <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Resumes_Observations conclusion].<br />
We will issue another Nugget when operations resume.<br />
<br />
== Annealing... ==<br />
<br />
Figure 2 shows the time history of the annealing procedure, from which a slow return to operating temperature is under way... to be continued....<br />
<br />
[[File:269new3.png|500px|thumb|center|Figure 2: <br />
Time history of the anneal, reflected in the temperature telemetered from the cold plate.]]<br />
<br />
The Sun has cooperated thus far; RHESSI has not missed any spectacular gamma-ray flares (yet).<br />
Figure 3 shows this cooperative spirit in terms of the GOES soft X-ray time history from the beginning of 2016.<br />
The most recent M-class flare was on 15 February.<br />
<br />
[[File:269new2.png|500px|thumb|center|Figure 3: <br />
GOES time history of the anneal, with the beginning of the anneal marked with a dashed line.<br />
Almost no flares at all!]]<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has article subject:: detector anneal| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI%27s_Anneal_AdventureRHESSI's Anneal Adventure2018-09-18T11:26:13Z<p>Schriste: /* Conclusions */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI's Anneal Adventure<br />
|number = 127<br />
|first_author = David Smith<br />
|second_author = Mark Lewis<br />
|publish_date = 2010 May 14<br />
|next_nugget = [[Awesome Stellar Flare Spectra]]<br />
|previous_nugget = [[History of Solar Oblateness]]<br />
}}<br />
<br />
==Introduction==<br />
<br />
RHESSI's germanium detectors periodically require maintenance in the form of [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 annealing].<br />
This just involves warming the detectors up so that their natural damage due to high-energy particles essentially fades away. <br />
The procedure is time-consuming and somewhat difficult, not to mention adventuresome, but the RHESSI team had planned to do its second anneal operation just about now in preparation for the major flares now beginning to happen as a part of [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_has_begun Solar Cycle 24].<br />
<br />
By an amazing coincidence, just as our planning cycle for the anneal operation was beginning, a totally unexpected glitch actually turned RHESSI off<br />
completely, including the cryocooler itself - an operational no-no, but of course the first step in an anneal operation since the detectors need to be heated, not cooled.<br />
RHESSI's cryocooler is a simplified and ruggedized version of a thermodynamic [http://en.wikipedia.org/wiki/Stirling_engine Stirling engine] and has the task of cooling the detectors - at the cost of some tens of watts of power - to below 100 K.<br />
The fact that this was totally unplanned made the operations a bit of an adventure, but the results seem to have been very satisfactory.<br />
This Nugget is just to keep RHESSI fans up to date on its status.<br />
This extends the information already available in the [http://sprg.ssl.berkeley.edu/RHESSI/recovery.txt RHESSI report] on the "anomaly", as provided to NASA.<br />
<br />
==Radiation Damage and Annealing==<br />
<br />
First, some background information.<br />
Energetic [http://imagine.gsfc.nasa.gov/docs/ask_astro/answers/970228a.html protons] and [http://en.wikipedia.org/wiki/Neutron_radiation neutrons] encountered in [http://marine.rutgers.edu/cool/education/class/paul/orbits2.html Earth orbit] penetrate<br />
RHESSI's [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 germanium detectors] and create small disordered regions in<br />
the germanium crystal lattice. <br />
When a solar X-ray or gamma ray interacts<br />
in a detector, it liberates a cloud of electron/hole pairs in the<br />
semiconductor that are swept to the detector's electrodes by the applied<br />
electric field, producing the current signal we detect with our<br />
electronics. <br />
But the disordered regions tend to become negatively charged<br />
and are effective [http://www.allaboutcircuits.com/vol_3/chpt_2/5.html hole] traps. <br />
By removing some of the charge moving<br />
through the crystal, they degrade the response of the detector; a<br />
monochromatic gamma-ray source, instead of looking like a nice, narrow<br />
Gaussian line in the spectrum, will grow a tail on the low-energy side<br />
that gets more and more severe as radiation damage increases. <br />
In about<br />
a minute, the trap will release the hole (too late to count, of course)<br />
and be ready to do further mischief. <br />
At the most extreme levels of<br />
radiation damage, the space density of negative charges from unfilled<br />
traps is so great that the resulting field can cancel the externally<br />
applied field, leaving parts of the detector volume with no field or a<br />
field pointing in the wrong direction. <br />
At this point, the detector has<br />
dead (passive) portions, most notably around the outside of the rear<br />
detector segments.<br />
<br />
[[Image:127fig1.jpg|600px|center|thumb|Figure 1: Results of the 2nd RHESSI anneal, for detector G9 only.<br />
The two panels show before/after (red/black) comparison background spectra for front and rear segments (see [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 earlier Nugget]).<br />
Note the dramatic improvement in resolution of the line features.]]<br />
<br />
For reasons that are not fully understood, heating the detectors to<br />
approximately 100 C will remove the tendency for most of the disordered<br />
regions to become negatively charged and function as hole traps. <br />
This<br />
temperature is not nearly high enough to re-crystallize the lattice, but<br />
this process is still referred to as "annealing".<br />
RHESSI's detectors were first annealed in November 2007, and the adventure<br />
reported here is the second time.<br />
Figure 1 shows the spectrum of<br />
background in a typical front and rear segment. <br />
Note the improvement<br />
both in the ability to resolve narrow background lines and in the total<br />
count rate (related to the active detector volume).<br />
<br />
==The Operations Adventure==<br />
<br />
On March 17 something happened on board RHESSI.<br />
We still do not know the exact details, but the leading idea is an unexpected undervoltage appeared <br />
on the main spacecraft [http://en.wikipedia.org/wiki/Electrical_bus power bus].<br />
This could have tripped an safety circuit that shut all systems down.<br />
The cryocoolers would normally continue functioning on battery power, but the safety trip<br />
also turned the battery heaters off.<br />
They froze, and no power was available during orbit night (this can be very, very cold in space).<br />
It did not help matters that this glitch appeared right after a series of <br />
[http://hessi.ssl.berkeley.edu/ground_systems/station.html Berkeley ground station] contact passes,<br />
when command control is normally available, so that the problem did not show up for half a day.<br />
<br />
However there was spacecraft power during the sunlit portions of the orbit, so eventually there was no<br />
problem about getting things sorted out.<br />
But in a case such as this, one is flying blind - the spacecraft is in an unplanned configuration, and so one <br />
must be very careful in recovering it.<br />
Considerable discussion of almost-forgotten spacecraft engineering ensued as the RHESSI team carefully <br />
nursed the spacecraft back into a "nominal" state (note that "anomalous" and "nominal" are NASA jargon<br />
basically meaning "bad" and "good" respectively).<br />
Everything recovered well enough so that the anneal could begin, and the silver lining of this story is that the<br />
RHESSI team thus did not have to spend a lot of time planning this delicate operation: it was forced on us.<br />
The "trips" have now been disabled, so that this cannot happen again, and the RHESSI team is working on<br />
replacing the positive features of this capability.<br />
<br />
The adventure here illustrates how tricky the operation of a delicate, remote spacecraft can be.<br />
For example, RHESSI pointing changed subtly as a result of the glitch.<br />
Such a change might be associated with something mechanical, such as a micrometeorite impact<br />
on one of the solar panels.<br />
Instead we think that it is just the re-initialization of the sequences of commands to the magnetic torquers <br />
that allow RHESSI to follow the <br />
[http://en.wikipedia.org/wiki/Galileo_Galilei apparent motion] of the Sun; these commands need to follow<br />
the spinning motion of the spacecraft and its geographical position, and are generated by software on board.<br />
<br />
==RHESSI images and spectra==<br />
<br />
[[Image:127fig2a.png|400px|center|thumb|Figure 2: Image of a small flare (contours) superposed on an earlier TRACE EUV image<br />
(courtesy Ryan Milligan).]]<br />
<br />
The real test of the success of the annealing operation (and the recovery from the glitch) is of course lies in<br />
the creation of images and spectra of real flares.<br />
The Sun has cooperated by having a few, even though full solar activity has not yet returned.<br />
The anneal (see Figure 1) clearly improved the spectral resolution of the detectors.<br />
Routine recalibration of the detector gains (pulse height vs. photon energy) has now also been accomplished,<br />
using the sharp background lines visible in that Figure. <br />
The bottom line is that RHESSI hard X-ray imaging works well again on an actual flare, as illustrated in Figure 2.<br />
Likewise we can make images with excellent spectral resolution again, as illustrated in Figure 3.<br />
This spectrum shows the Fe feature at 6.7 keV, a good diagnostic of the thermal plasma; there is also a hard X-ray<br />
tail that measures the non-thermal electron content of the event.<br />
The spectrum (Figure 3) and the image (Figure 2) are from different small flares.<br />
<br />
<br />
[[Image:127fig3.png|500px|center|thumb|Figure 3: Spectral fits for two small flares, one before and one after the successful anneal operation<br />
(courtesy Amir Caspi).<br />
Note the presence of the Fe line feature at 6.7 keV, confirming the high spectral resolution now being achieved again.]]<br />
<br />
==Conclusions==<br />
<br />
RHESSI has brilliantly survived its second anneal, even though it occurred somewhat anomalously.<br />
The best results will come from the flare observations to be obtained soon, but we know enough already to be very pleased indeed.<br />
<br />
* The overall radiation damage state in the rear segments is about as it was in late 2004 or early 2005.<br />
* There is now no appreciable degradation of the front segments.<br />
* Detector 7 is working nominally, with a 3 keV threshold, which it has never done before, not even on the ground! The anneal may have purged contaminants off the detector's surface.<br />
* Detector 2, always the problem child, also has a 3 keV threshold, although poorer resolution than the other detectors.<br />
* All detectors are holding > 3000 V in high voltage.<br />
* The new detector gains have been characterized.<br />
* RHESSI's rate of detection of Terrestrial Gamma-ray Flashes -- one of its most interesting non-solar science products --<br />
has returned to pre-damage levels.<br />
<br />
Overall, RHESSI looks far healthier now than before our involuntary second anneal, and even better than just after the first (planned) one.<br />
We look forward now to stalwart service by RHESSI, the only instrument capable of imaging the fundamentally important hard X-ray<br />
and &gamma;-ray emissions of solar flares.<br />
With the new [http://sdo.gsfc.nasa.gov/ Solar Dynamics Observatory] successfully in orbit, in fact, RHESSI now has a comprehensive "context" observatory that should make every observation count.<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has article subject:: detector anneal| ]]<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has observation by:: TRACE| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_is_Annealing_NowRHESSI is Annealing Now2018-09-18T02:41:31Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI is Annealing Now<br />
|first_author = Albert Shih <br />
|second_author = Martin Fivian <br />
|publish_date = July 7, 2014<br />
|previous_nugget = [[Mysteries of Flare/CME Initiation]]<br />
|next_nugget = [[The Redistribution of Nonthermal Electron Energy]] <br />
|number = 229<br />
}}<br />
<br />
== Introduction ==<br />
<br />
RHESSI's detectors are high-resolution <br />
[https://en.wikipedia.org/wiki/Gamma_spectroscopy germanium spectrometers], widely used <br />
in many situations that require gamma-ray spectroscopy, for example in the measurement<br />
of emission lines from <br />
[http://hyperphysics.phy-astr.gsu.edu/hbase/nuclear/shell.html excited nuclear states].<br />
Solar flares have these, and they provide the most direct way to learn about the <br />
essential process of particle acceleration in these eruptions.<br />
<br />
Germanium detectors in space accumulate<br />
[https://en.wikipedia.org/wiki/Radiation_damage radiation damage]<br />
in the hostile environment above Earth's protective atmosphere.<br />
This results in reduced resolution and calibration errors, as a result of crystal defects that appear<br />
in the germanium.<br />
The time scale for this degradation is a few years, but it shortens as time goes on.<br />
Luckily, it turns out that simply heating the detectors up to about 100 degrees Centigrade for a week to ten days can reset these defects<br />
via an<br />
[http://www.thefreedictionary.com/anneal annealing] <br />
process.<br />
<br />
We do not anneal the RHESSI detectors very often, since each anneal requires about six weeks for the very <br />
carefully done heating and cooling processes.<br />
An [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 earlier Nugget] describes<br />
the result of the first anneal, in 2008, and also details some of RHESSI's special properties.<br />
<br />
== The current anneal ==<br />
<br />
The annealing operation now under way is the fourth, and we anticipate that it will restore some of the<br />
degradation accumulated since the previous anneal in 2012.<br />
Figure 1 shows the time histories of detector ("cold plate") temperatures for the four processes; at the<br />
time of writing the process is about halfway through its "hot" (about 100 C) phase and the slow<br />
process of returning to a correct operating temperature will commence in mid-July.<br />
<br />
[[Image:229f1.png|500px|thumb|center|'''Figure 1''': <br />
Time history of RHESSI temperatures during the three prior annealings, plus the beginning of the present one.<br />
]]<br />
<br />
Normal operations will begin in mid-August and in the meanwhile, the RHESSI team was hoping that the Sun <br />
would not produce any major flares.<br />
But our hopes were dashed; almost as soon as the anneal began, a quite remarkable event occurred, as <br />
shown in Figure 2 (see also <br />
[http://sdowww.lmsal.com/sdomedia/ssw/media/ssw/ssw_client/data/ssw_service_140626_215545_49675/www/ Sam Freeland's coverage]<br />
of this event, which was a coronal disturbance originating from a highly occulted flare).<br />
The occurrence of this event during RHESSI's down time just emphasizes the importance of our unique<br />
instrument. Apart from the [http://hesperia.gsfc.nasa.gov/fermi_solar/ Gamma-ray Burst Monitor on the Fermi spacecraft] that provides X-ray and gamma-ray spectroscopy above ~10 keV with relatively modest energy resolution but no imaging, RHESSI will be the only source of high-resolution X-ray imaging spectroscopy prior to the launch of the <br />
[http://stix.i4ds.ch/mission/ STIX]<br />
instrument on<br />
[http://sci.esa.int/solar-orbiter/ Solar Orbiter] several years from now.<br />
<br />
[[Image:229f2.png|600px|thumb|center|'''Figure 2''': <br />
A remarkable "gradual rise and fall" event, overlapping two days of <br />
[http://sprg.ssl.berkeley.edu/~tohban/browser/ RHESSI Browser] standard plots, now annotated to show the annealing operation in progress.<br />
This event originated on the invisible hemisphere of the Sun, and would likely have produced interesting coronal hard X-ray sources for RHESSI to study (Ref. [1]).<br />
]]<br />
<br />
== What do we expect? ==<br />
<br />
When the anneal has been completed, we will write another Nugget on the results.<br />
At a minimum we expect that the energy resolution and calibration of most of RHESSI's nine independent<br />
detectors will improve drastically.<br />
This should enable RHESSI to observe the remainder of the present solar maximum with high-quality <br />
hard X-ray imaging spectroscopy, as it has done since <br />
[http://hesperia.gsfc.nasa.gov/hessi/hessi_launch.htm launch] in February, 2002.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2008A%26ARv..16..155K "Hard X-ray emission from the solar corona"]<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has article subject:: detector anneal| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/The_RHESSI_Flare_CatalogThe RHESSI Flare Catalog2018-09-18T02:40:17Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = The RHESSI Flare Catalog<br />
|number = 183<br />
|first_author = Jim McTiernan <br />
|second_author = Hugh Hudson<br />
|publish_date = 29 August 2012<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::184]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::182]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
RHESSI recently celebrated its [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI%27s_Tenth_Anniversary 10th anniversary], and the length of its <br />
database now approaches the length of an 11-year [http://www.swpc.noaa.gov/SolarCycle/ solar cycle], albeit<br />
from maximum to maximum (2002-2012 at the time of writing).<br />
<br />
The objective of this Nugget is to advertise the RHESSI flare catalog, a<br />
unique database of hard X-ray flare occurrence - the second, following the<br />
[http://solar.physics.montana.edu/mckenzie/HXTflarelist.html Yohkoh/HXT catalog] (1991-2001), to include flare positions and<br />
spectral information.<br />
We discuss these as an application of the catalog in the next section<br />
of this Nugget.<br />
<br />
== The catalog ==<br />
<br />
It is very easy to access the catalog: from any computer using [http://www.lmsal.com/solarsoft/ SolarSoft] and having<br />
the RHESSI software, one simply enters<br />
<br />
IDL> flares = hsi_read_flarelist()<br />
<br />
and<br />
<br />
IDL> help, flares, /structure<br />
<br />
to see the contents (including [http://www.swpc.noaa.gov/today.html GOES] flare classifications and NOAA [http://www.solarmonitor.org/ active-region] identifications). <br />
There is also a GUI for selecting flares which you get by typing <br />
<br />
IDL> hsi_flarecat, list=flares, /struct<br />
<br />
See the [http://hesperia.gsfc.nasa.gov/rhessi2/ RHESSI Web pages] for much more explanation of the data and<br />
the catalog database. For a detailed explanation of the creation of the flare catalog, <br />
see [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Hsi_flare_list_fill:_How_the_RHESSI_Flare_List_is_Generated this Wiki article.] <br />
The RHESSI [http://sprg.ssl.berkeley.edu/~tohban/browser/ Browser], as described in an earlier [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=20 Nugget],<br />
shows simple quick-look light curves and images for each flare in the <br />
catalog. <br />
As of 20 August 2012 the catalog held 63,984 entries, for a mean flare<br />
rate of more than 16 entries per day.<br />
For comparison the event list, for the same time interval, <br />
contained some 16,206 events, even with almost continuous coverage<br />
(unlike RHESSI's, which has appreciable gaps due to its low Earth<br />
orbit).<br />
The RHESSI catalog is therefore is a bit deeper, though we note that some flares trigger the catalog criteria multiple times.<br />
Its other advantages include locations and spectra, and also<br />
(because RHESSI observes hard X-ray events, which have shorter durations) <br />
less background confusion for fainter events.<br />
GOES notoriously under-counts flare even at the C class during solar maxima,<br />
simply because the events pile up upon each other.<br />
<br />
== An application ==<br />
<br />
One easy application of the RHESSI catalog has to do with flare<br />
occcurrence patterns.<br />
Earlier Nuggets ([http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=52 1], [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=81 2]) have already described the systematic appearance<br />
of the flare locations thus derived.<br />
Figure 1 shows two graphic representations of the catalog positions:<br />
a simple [x,y] plot, with each point denoting one entry; and a plot of<br />
radial distances (in both cases, the measurement refers to displacements<br />
in arc sec from Sun Center).<br />
<br />
[[File:183f1.png|thumb|left|400px|Figure 1:<br />
Raw map of catalog flares viewed against the solar disk (the two colors show its maximum and minimum angular extent).<br />
This map includes artifactual errors that can easily be screened out.<br />
]]<br />
<br />
[[File:183f2.png|thumb|right|400px|Figure 2:<br />
Distribution of flare radial locations at the extreme limb, now converted to units of solar radii.<br />
The vertical dotted line shows the theoretical height of the X-ray limb, as determined via model atmospheres and the cross-section<br />
for [http://en.wikipedia.org/wiki/Compton_scattering Compton scattering].<br />
]]<br />
<br />
The plots are informative.<br />
On the left one can see some of the artifacts present in the catalog;<br />
these include events at high latitudes.<br />
Ref. [1] studied an earlier version of this distribution and concluded<br />
that all of these were artifactual, and that all of the RHESSI flare<br />
and flare-like events in fact come from active regions, which do not<br />
occur above latitudes of about 40 degrees.<br />
Note that the horizontal streak of points at [X,Y] = [200-500,700]<br />
is in fact real, and these record-setting events include an M-class<br />
flare that we may describe in an future Nugget.<br />
<br />
The plot on the right (radial distances) ignores all of the artifacts.<br />
This should be a lesson in data analysis; one can see the artifacts in the<br />
(X,Y) map clearly enough, and a correct analysis procedure would involve<br />
screening the catalog entries to remove these.<br />
A main source of problems here comes from the (rare) RHESSI [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=30 offpoints],<br />
and other occasions where the pointing coordinates are unreliable.<br />
<br />
Bearing that in mind, Figure 2 (right) matches expectations: the flare hard <br />
X-ray sources form a cloud in the solar corona, and their distribution must<br />
have the appearance of an optically-thin shell.<br />
This would have a jump by precisely a factor of two at the X-ray limb <br />
(shown as estimated theoretically as a dotted line in the Figure).<br />
Any events located above this point would lie at high altitudes<br />
(see the [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=41 Nugget] on coronal hard X-ray sources).<br />
Of course, angular resolution is an issue here: the FWHM of the RHESSI<br />
[collimators] used for the catalog is only about 0.01 R<sub>Sun</sub>, and<br />
so there are comparable uncertainties in a distribution made in this way.<br />
<br />
== Caveats regarding use of the flare catalog ==<br />
<br />
1.) As mentioned above, position information can be inaccurate, typically due to <br />
difficulties with the spacecraft pointing solution, and high latitude positions <br />
in particular are not reliable.<br />
<br />
2.) The RHESSI flare list over counts flares, especially when handling long-duration events. In particular, data gaps of longer than 5 minutes, caused by spacecraft night, SAA passage, or missing data will cause a split of one flare into more than one events. Also events that some observers might call single flares with multiple peaks are often split into separate flares; this is a consequence of the fact that flares in the 6 to 12 keV range overlap in time. The process used to split flares is documented in [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Hsi_flare_list_fill:_How_the_RHESSI_Flare_List_is_Generated this Wiki article.]<br />
<br />
3.) GOES and NOAA information often lags behind the RHESSI processing, so data regarding active regions and GOES class is often missing. This is usually picked up during later reprocessing, but users are encouraged to regard these quantities as informational, and not attempt quantitative analysis without careful testing.<br />
<br />
4.) The RHESSI flare list is currently undergoing re-processing, to fix a problem with calculation of X-ray photon flux. A new version should appear in late 2012.<br />
<br />
== Conclusions ==<br />
<br />
We have briefly described the RHESSI flare catalog and sketched out one<br />
application (the determination of hard X-ray source heights).<br />
At first look the catalog data are only a crude approximation, with <br />
low resolution and minimal sampling.<br />
And yet the histogram shown in Figure 1 (right) can be made with a binning<br />
of 0.5 arc s, a "super-resolution" factor of 10 or so.<br />
Such is the potential power of a very large database, but of course the <br />
harder one pushes the precision of such a distribution, the greater the<br />
influence of systematic errors.<br />
We definitely encourage users to explore the catalog, consider its systematic<br />
effects, and study large volumes of RHESSI data to learn precise things about<br />
flare occurrence.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2008ApJ...677..704H RHESSI Microflare Statistics. II. X-Ray Imaging, Spectroscopy, and Energy Distributions]<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has article subject:: education| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI%27s_Anneal_AdventureRHESSI's Anneal Adventure2018-09-18T02:39:27Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI's Anneal Adventure<br />
|number = 127<br />
|first_author = David Smith<br />
|second_author = Mark Lewis<br />
|publish_date = 2010 May 14<br />
|next_nugget = [[Awesome Stellar Flare Spectra]]<br />
|previous_nugget = [[History of Solar Oblateness]]<br />
}}<br />
<br />
==Introduction==<br />
<br />
RHESSI's germanium detectors periodically require maintenance in the form of [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 annealing].<br />
This just involves warming the detectors up so that their natural damage due to high-energy particles essentially fades away. <br />
The procedure is time-consuming and somewhat difficult, not to mention adventuresome, but the RHESSI team had planned to do its second anneal operation just about now in preparation for the major flares now beginning to happen as a part of [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_has_begun Solar Cycle 24].<br />
<br />
By an amazing coincidence, just as our planning cycle for the anneal operation was beginning, a totally unexpected glitch actually turned RHESSI off<br />
completely, including the cryocooler itself - an operational no-no, but of course the first step in an anneal operation since the detectors need to be heated, not cooled.<br />
RHESSI's cryocooler is a simplified and ruggedized version of a thermodynamic [http://en.wikipedia.org/wiki/Stirling_engine Stirling engine] and has the task of cooling the detectors - at the cost of some tens of watts of power - to below 100 K.<br />
The fact that this was totally unplanned made the operations a bit of an adventure, but the results seem to have been very satisfactory.<br />
This Nugget is just to keep RHESSI fans up to date on its status.<br />
This extends the information already available in the [http://sprg.ssl.berkeley.edu/RHESSI/recovery.txt RHESSI report] on the "anomaly", as provided to NASA.<br />
<br />
==Radiation Damage and Annealing==<br />
<br />
First, some background information.<br />
Energetic [http://imagine.gsfc.nasa.gov/docs/ask_astro/answers/970228a.html protons] and [http://en.wikipedia.org/wiki/Neutron_radiation neutrons] encountered in [http://marine.rutgers.edu/cool/education/class/paul/orbits2.html Earth orbit] penetrate<br />
RHESSI's [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 germanium detectors] and create small disordered regions in<br />
the germanium crystal lattice. <br />
When a solar X-ray or gamma ray interacts<br />
in a detector, it liberates a cloud of electron/hole pairs in the<br />
semiconductor that are swept to the detector's electrodes by the applied<br />
electric field, producing the current signal we detect with our<br />
electronics. <br />
But the disordered regions tend to become negatively charged<br />
and are effective [http://www.allaboutcircuits.com/vol_3/chpt_2/5.html hole] traps. <br />
By removing some of the charge moving<br />
through the crystal, they degrade the response of the detector; a<br />
monochromatic gamma-ray source, instead of looking like a nice, narrow<br />
Gaussian line in the spectrum, will grow a tail on the low-energy side<br />
that gets more and more severe as radiation damage increases. <br />
In about<br />
a minute, the trap will release the hole (too late to count, of course)<br />
and be ready to do further mischief. <br />
At the most extreme levels of<br />
radiation damage, the space density of negative charges from unfilled<br />
traps is so great that the resulting field can cancel the externally<br />
applied field, leaving parts of the detector volume with no field or a<br />
field pointing in the wrong direction. <br />
At this point, the detector has<br />
dead (passive) portions, most notably around the outside of the rear<br />
detector segments.<br />
<br />
[[Image:127fig1.jpg|600px|center|thumb|Figure 1: Results of the 2nd RHESSI anneal, for detector G9 only.<br />
The two panels show before/after (red/black) comparison background spectra for front and rear segments (see [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 earlier Nugget]).<br />
Note the dramatic improvement in resolution of the line features.]]<br />
<br />
For reasons that are not fully understood, heating the detectors to<br />
approximately 100 C will remove the tendency for most of the disordered<br />
regions to become negatively charged and function as hole traps. <br />
This<br />
temperature is not nearly high enough to re-crystallize the lattice, but<br />
this process is still referred to as "annealing".<br />
RHESSI's detectors were first annealed in November 2007, and the adventure<br />
reported here is the second time.<br />
Figure 1 shows the spectrum of<br />
background in a typical front and rear segment. <br />
Note the improvement<br />
both in the ability to resolve narrow background lines and in the total<br />
count rate (related to the active detector volume).<br />
<br />
==The Operations Adventure==<br />
<br />
On March 17 something happened on board RHESSI.<br />
We still do not know the exact details, but the leading idea is an unexpected undervoltage appeared <br />
on the main spacecraft [http://en.wikipedia.org/wiki/Electrical_bus power bus].<br />
This could have tripped an safety circuit that shut all systems down.<br />
The cryocoolers would normally continue functioning on battery power, but the safety trip<br />
also turned the battery heaters off.<br />
They froze, and no power was available during orbit night (this can be very, very cold in space).<br />
It did not help matters that this glitch appeared right after a series of <br />
[http://hessi.ssl.berkeley.edu/ground_systems/station.html Berkeley ground station] contact passes,<br />
when command control is normally available, so that the problem did not show up for half a day.<br />
<br />
However there was spacecraft power during the sunlit portions of the orbit, so eventually there was no<br />
problem about getting things sorted out.<br />
But in a case such as this, one is flying blind - the spacecraft is in an unplanned configuration, and so one <br />
must be very careful in recovering it.<br />
Considerable discussion of almost-forgotten spacecraft engineering ensued as the RHESSI team carefully <br />
nursed the spacecraft back into a "nominal" state (note that "anomalous" and "nominal" are NASA jargon<br />
basically meaning "bad" and "good" respectively).<br />
Everything recovered well enough so that the anneal could begin, and the silver lining of this story is that the<br />
RHESSI team thus did not have to spend a lot of time planning this delicate operation: it was forced on us.<br />
The "trips" have now been disabled, so that this cannot happen again, and the RHESSI team is working on<br />
replacing the positive features of this capability.<br />
<br />
The adventure here illustrates how tricky the operation of a delicate, remote spacecraft can be.<br />
For example, RHESSI pointing changed subtly as a result of the glitch.<br />
Such a change might be associated with something mechanical, such as a micrometeorite impact<br />
on one of the solar panels.<br />
Instead we think that it is just the re-initialization of the sequences of commands to the magnetic torquers <br />
that allow RHESSI to follow the <br />
[http://en.wikipedia.org/wiki/Galileo_Galilei apparent motion] of the Sun; these commands need to follow<br />
the spinning motion of the spacecraft and its geographical position, and are generated by software on board.<br />
<br />
==RHESSI images and spectra==<br />
<br />
[[Image:127fig2a.png|400px|center|thumb|Figure 2: Image of a small flare (contours) superposed on an earlier TRACE EUV image<br />
(courtesy Ryan Milligan).]]<br />
<br />
The real test of the success of the annealing operation (and the recovery from the glitch) is of course lies in<br />
the creation of images and spectra of real flares.<br />
The Sun has cooperated by having a few, even though full solar activity has not yet returned.<br />
The anneal (see Figure 1) clearly improved the spectral resolution of the detectors.<br />
Routine recalibration of the detector gains (pulse height vs. photon energy) has now also been accomplished,<br />
using the sharp background lines visible in that Figure. <br />
The bottom line is that RHESSI hard X-ray imaging works well again on an actual flare, as illustrated in Figure 2.<br />
Likewise we can make images with excellent spectral resolution again, as illustrated in Figure 3.<br />
This spectrum shows the Fe feature at 6.7 keV, a good diagnostic of the thermal plasma; there is also a hard X-ray<br />
tail that measures the non-thermal electron content of the event.<br />
The spectrum (Figure 3) and the image (Figure 2) are from different small flares.<br />
<br />
<br />
[[Image:127fig3.png|500px|center|thumb|Figure 3: Spectral fits for two small flares, one before and one after the successful anneal operation<br />
(courtesy Amir Caspi).<br />
Note the presence of the Fe line feature at 6.7 keV, confirming the high spectral resolution now being achieved again.]]<br />
<br />
==Conclusions==<br />
<br />
RHESSI has brilliantly survived its second anneal, even though it occurred somewhat anomalously.<br />
The best results will come from the flare observations to be obtained soon, but we know enough already to be very pleased indeed.<br />
<br />
* The overall radiation damage state in the rear segments is about as it was in late 2004 or early 2005.<br />
* There is now no appreciable degradation of the front segments.<br />
* Detector 7 is working nominally, with a 3 keV threshold, which it has never done before, not even on the ground! The anneal may have purged contaminants off the detector's surface.<br />
* Detector 2, always the problem child, also has a 3 keV threshold, although poorer resolution than the other detectors.<br />
* All detectors are holding > 3000 V in high voltage.<br />
* The new detector gains have been characterized.<br />
* RHESSI's rate of detection of Terrestrial Gamma-ray Flashes -- one of its most interesting non-solar science products --<br />
has returned to pre-damage levels.<br />
<br />
Overall, RHESSI looks far healthier now than before our involuntary second anneal, and even better than just after the first (planned) one.<br />
We look forward now to stalwart service by RHESSI, the only instrument capable of imaging the fundamentally important hard X-ray<br />
and &gamma;-ray emissions of solar flares.<br />
With the new [http://sdo.gsfc.nasa.gov/ Solar Dynamics Observatory] successfully in orbit, in fact, RHESSI now has a comprehensive "context" observatory that should make every observation count.<br />
<br />
[[Has article subject:: RHESSI| ]]<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has observation by:: TRACE| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/History_of_Solar_OblatenessHistory of Solar Oblateness2018-09-18T02:36:55Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = History of Solar Oblateness<br />
|number = 126<br />
|first_author = Hugh Hudson<br />
|second_author = Jean-Pierre Rozelot<br />
|publish_date = 2010 April 26<br />
|next_nugget = [[RHESSI's Anneal Adventure]]<br />
|previous_nugget = [[An Alternative View of the Masuda Flare]]<br />
}}<br />
<br />
== Introduction ==<br />
<br />
Exact measurements of the shape of the Sun have a history extending well back into the 19th century (for full details, see reference [1]), and RHESSI is playing a small role in this continuing history (see two earlier nuggets<br />
([http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=19 a], [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=85 b]) thanks to the three small, simple optical telescopes it uses for solar aspect sensing.<br />
These observations are [http://en.wikipedia.org/wiki/The_Three_Princes_of_Serendip serendipitous], as are those of [http://sohowww.nascom.nasa.gov/ SOHO/MDI] and soon those of [http://hmi.stanford.edu/ SDO/HMI].<br />
Other dedicated space observatories are being actually <i>designed</i> for solar global observations, and we expect that these will make definitive measurements over the next few years.<br />
In related fields we note the tremendous success of the photometry/[http://en.wikipedia.org/wiki/Asteroseismology asteroseismology] missions [http://www.astro.ubc.ca/MOST/ MOST], [http://132.149.11.177/COROT/GP_satellite.htm Corot], and [http://kepler.nasa.gov/ Kepler]. <br />
<br />
The purpose of this Nugget is to show off some of the historical overview from [1], and to remind Nugget readers of the basic physics of the oblateness measurement. <br />
RHESSI has just completed its observations through the remarkable [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_has_begun recent solar minimum] and we see no particular reason why its data should not continue well into the operational lifetimes of SDO and Picard.<br />
<br />
== Why is the oblateness interesting? ==<br />
<br />
[http://en.wikipedia.org/wiki/Astrometry Astrometry] is one of the classical branches of astronomy, and precise determinations of the shape of the Sun make use of highly specialized techniques with similar fundamental problems of measurement.<br />
The oblateness of the Sun is normally defined as the (normalized) difference between equatorial and polar radii, and so it is a differential measurement as opposed to the absolute determination of the radius (or diameter), which is obviously much harder.<br />
The RHESSI, SOHO, and now SDO measurements cannot be regarded as absolute and are restricted to shape alone, but we look forward to optimized and more direct measurements that can precisely determine the radius as well as the shape.<br />
<br />
[[Image:126fig1.jpg|right|thumb|250px|Figure 1: Cross-section of the Sun through its rotation axis, showing its basically circular shape (red) with a small distortion induced by the rotation (blue).<br />
The radial scale has been magnified enormously as indicated.<br />
The points show actual RHESSI data as described in the text.]]<br />
<br />
The shape of the Sun, or any star, reflects what is going on inside it.<br />
Most trivially it is rotating, and so an equatorial bulge should appear.<br />
This would be the main source of non-circularity sketched in Figure 1.<br />
The Sun rotates slowly, and so this bulge is small (by one prediction, 7.98 mas, where the common unit mas is a milli-arc-sec, or closely 1 part per million of the<br />
full radius).<br />
More rapidly rotating stars can have substantial bulge, which leads to the interesting effect known as the [http://en.wikipedia.org/wiki/Von_Zeipel_theorem von Zeipel theorem] which implies that there is also a substantial dimming of the surface brightness at the equator of such a star.<br />
In the case of the Sun both rotational ellipticity and rotational dimming are tiny and therefore a wonderful challenge for observers.<br />
<br />
Many other potential mechanisms could change the shape of the Sun.<br />
The surface is known to [http://www.aip.de/~mku/difrot/difrot.html rotate differentially] in the sense that the equator region has an angular velocity greater than the polar regions.<br />
There must be internal differential rotation as well, although it is increasingly well constrained by [http://soi.stanford.edu/results/heliowhat.html helioseismology], and what we see at the surface may conceal mass internal mass distributions that can affect the shape.<br />
Such effects are related to [http://en.wikipedia.org/wiki/Tests_of_general_relativity tests of relativity theory], from which Einstein famously explained the anomalous perihelion precession of Mercury.<br />
The observation of this precession was an amazing coup of 19th-century astrometry.<br />
Finally there are links to the [http://en.wikipedia.org/wiki/Solar_cycle solar cycle]; for example a north-south [http://www.noao.edu/noao/staff/irenegh/meridional_web/meridional.html meridional circulation] is required to explain the 11-year cycle - a hot research topic nowadays because of the<br />
recent [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_-_don%27t_panic_yet%21 anomalous minimum].<br />
This could well have detectable effects at the surface of the Sun other than the subtle motions of magnetic elements that trace out the flow associated with the cycle.<br />
<br />
== A history of the measurements ==<br />
<br />
[[Image:126fig2.jpg|left|thumb|300px|Figure 2: Early measurements of the solar oblateness, using both photographic plates and [http://en.wikipedia.org/wiki/Heliometer heliometers]. <br />
These early researchers were interested in the relationship between the oblateness and the phase of the solar [http://www.solarcycle24.com/ sunspot cycle], a question remaining open in spite of the great improvement in measurement accuracy since the 19th century.]]<br />
<br />
Early measurements of solar oblateness used the classical techniques of astrometry such as the measurement carefully exposed photographic plates, or visual observations through transit measurements or via [http://en.wikipedia.org/wiki/Heliometer heliometers] specifically designed for this purpose.<br />
The precision of such measurements (see Figure 2), although well below an arc second, is far from what is known now to be necessary, but that did not keep the observers from speculating about the apparent variations (all due to measurement uncertainty) that they detected - the solid line in Figure 2 compares the observations with the sunspot number.<br />
The [http://en.wikipedia.org/wiki/Astrolabe astrolabe] can also be adapted for measurements of the solar diameter, and can be instrumented with modern detectors as well.<br />
The modern era of solar oblateness measurements can be said to have begun with [http://en.wikipedia.org/wiki/Robert_H._Dicke Dicke], who wished to make measurements sufficiently accurate to test rival theories of [http://en.wikipedia.org/wiki/General_relativity general relativity].<br />
Dicke's ground-based telescope contained several innovations designed to make the measurement differential in nature and as immune as possible to [http://www.schoolsobservatory.org.uk/astro/tels/seeing.shtml "seeing"] conditions.<br />
The device largely succeeded in this, especially after it was moved from Princeton, New Jersey, to a better site in California.<br />
It also inspired other instrumental development.<br />
<br />
Systematic observations of solar oblateness now continue from ground-based observatories, and the measurement errors have dropped to the order of one mas (milli arc sec, or about a part per million).<br />
The new era, as mentioned, began with the observations of Dicke & Goldenberg as reported in [2], but these data were quickly superseded by superior observations, also ground-based.<br />
Figure 2 shows how the error bars dropped during this intermediate period of observation, roughly speaking from 1960 to 1990 (see [1] for details of the individual data).<br />
<br />
[[Image:126fig3.jpg|left|thumb|400px|Figure 3: The intermediate period of specialized, modern ground-based observations of solar oblateness.<br />
The first point on this graph was the revolutionary observation of Dicke & Goldenberg, as reported in [2].]]<br />
<br />
[[Image:126fig4.jpg|right|thumb|400px|Figure 4: The modern data, including the single point thus far published from the RHESSI data.<br />
The black lozenges show the continuing ground-based observations from [http://en.wikipedia.org/wiki/Pic_du_Midi_de_Bigorre Pic du Midi] (see also the Pic du Midi [http://ljr.bagn.obs-mip.fr/histo/histo.html home page]).<br />
The green points show the recent solar maximum; please refer to [1] for fuller details.]]<br />
<br />
It has now become possible to do observations from balloon-borne platforms and from actual satellites in space.<br />
Because the measurements require such great precision, general-purpose solar telescopes are not very useful for oblateness measurements in general, although any full-disk image can be analyzed for its astrometric content.<br />
Thus the observations reported from [http://soi.stanford.edu/ SOHO/MDI] and soon, we anticipate, from [http://hmi.stanford.edu/ SDO/HMI], are lineal descendants of the astrometric plate reductions shown in Figure 1.<br />
Nowadays of course the plates have been replaced by [http://nobelprize.org/nobel_prizes/physics/laureates/2009/press.html CCDs].<br />
<br />
The RHESSI observations [3] presently hold the record in terms of precision, with an oblateness uncerainty of only 0.14 mas. <br />
This has been possible because of a fortunate and accidental imitation of one of Dicke's great ideas: to have the telescope rotate rapidly. <br />
By this artifice many systematic errors become more tractable, since the rotation allows the observer to study deviations from a circular shape almost directly. <br />
The lowest-order deviation is the oblateness term, but as mentioned in the Introduction we fully expect that more precise measurements will reveal higher-order shape terms.<br />
Figure 4 shows the single RHESSI data point, but it is hard to spot at first because of its small error bars.<br />
Eventually we hope to have a full time series of comparable data, extending at least from 2002 to the time of writing (2010).<br />
Other notable data from the current epoch include a continuing series of high-precision measurements from<br />
[http://en.wikipedia.org/wiki/Pic_du_Midi_de_Bigorre Pic du Midi], shown as lozenges with error flags in Figure 4.<br />
The green points in Figure 4 show solar activity, but in terms of the [http://www.scholarpedia.org/article/Solar_activity facular index] rather than sunspot number.<br />
<br />
The RHESSI data, as plotted, show an apparent oblateness greater than 10 mas.<br />
As described in [3], this appears to be the result of solar [http://www.scholarpedia.org/article/Solar_activity magnetic activity] in the form of faculae and enhanced network; the elimination of this contaminating signal brought the apparent oblateness down to the level predicted for uniform rotation.<br />
This result needs confirmation, which we hope to have soon both from RHESSI and from other space missions as well as the ground-based observations.<br />
<br />
== Conclusions ==<br />
<br />
The measurement of solar oblateness has a rich history and a some profound implications touching on relativity theory, for example. <br />
RHESSI is making its contributions to this history [3].<br />
The shape of the Sun is one of the ways we have now for peering into its interior and learning empirically about flows and motions there that would otherwise only be guessed from theoretical considerations.<br />
The paramount of these techniques is of course helioseismology, a technique now being extended to the stars.<br />
<br />
== References ==<br />
<br />
[1] [http://www.sciencedirect.com/science?_ob=ArticleURL&_udi=B6VHB-4YJ4NKR-4&_user=121723&_coverDate=03%2F06%2F2010&_alid=1307462546&_rdoc=2&_fmt=high&_orig=search&_cdi=6062&_sort=r&_docanchor=&view=c&_ct=5&_acct=C000009999&_version=1&_urlVersion=0&_userid=121723&md5=0f904d940638a417b1641698fd30b82f A brief history of the solar oblateness. A review]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/1967PhRvL..18..313D Solar oblateness and general relativity]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2008Sci...322..560F A large excess in apparent solar oblateness due to surface magnetism]<br />
<br />
[[Has article subject:: history| ]]<br />
[[Has article subject:: solar oblateness| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/History_of_Solar_OblatenessHistory of Solar Oblateness2018-09-18T02:36:31Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = History of Solar Oblateness<br />
|number = 126<br />
|first_author = Hugh Hudson<br />
|second_author = Jean-Pierre Rozelot<br />
|publish_date = 2010 April 26<br />
|next_nugget = [[RHESSI's Anneal Adventure]]<br />
|previous_nugget = [[An Alternative View of the Masuda Flare]]<br />
}}<br />
<br />
== Introduction ==<br />
<br />
Exact measurements of the shape of the Sun have a history extending well back into the 19th century (for full details, see reference [1]), and RHESSI is playing a small role in this continuing history (see two earlier nuggets<br />
([http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=19 a], [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=85 b]) thanks to the three small, simple optical telescopes it uses for solar aspect sensing.<br />
These observations are [http://en.wikipedia.org/wiki/The_Three_Princes_of_Serendip serendipitous], as are those of [http://sohowww.nascom.nasa.gov/ SOHO/MDI] and soon those of [http://hmi.stanford.edu/ SDO/HMI].<br />
Other dedicated space observatories are being actually <i>designed</i> for solar global observations, and we expect that these will make definitive measurements over the next few years.<br />
In related fields we note the tremendous success of the photometry/[http://en.wikipedia.org/wiki/Asteroseismology asteroseismology] missions [http://www.astro.ubc.ca/MOST/ MOST], [http://132.149.11.177/COROT/GP_satellite.htm Corot], and [http://kepler.nasa.gov/ Kepler]. <br />
<br />
The purpose of this Nugget is to show off some of the historical overview from [1], and to remind Nugget readers of the basic physics of the oblateness measurement. <br />
RHESSI has just completed its observations through the remarkable [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_has_begun recent solar minimum] and we see no particular reason why its data should not continue well into the operational lifetimes of SDO and Picard.<br />
<br />
== Why is the oblateness interesting? ==<br />
<br />
[http://en.wikipedia.org/wiki/Astrometry Astrometry] is one of the classical branches of astronomy, and precise determinations of the shape of the Sun make use of highly specialized techniques with similar fundamental problems of measurement.<br />
The oblateness of the Sun is normally defined as the (normalized) difference between equatorial and polar radii, and so it is a differential measurement as opposed to the absolute determination of the radius (or diameter), which is obviously much harder.<br />
The RHESSI, SOHO, and now SDO measurements cannot be regarded as absolute and are restricted to shape alone, but we look forward to optimized and more direct measurements that can precisely determine the radius as well as the shape.<br />
<br />
[[Image:126fig1.jpg|right|thumb|250px|Figure 1: Cross-section of the Sun through its rotation axis, showing its basically circular shape (red) with a small distortion induced by the rotation (blue).<br />
The radial scale has been magnified enormously as indicated.<br />
The points show actual RHESSI data as described in the text.]]<br />
<br />
The shape of the Sun, or any star, reflects what is going on inside it.<br />
Most trivially it is rotating, and so an equatorial bulge should appear.<br />
This would be the main source of non-circularity sketched in Figure 1.<br />
The Sun rotates slowly, and so this bulge is small (by one prediction, 7.98 mas, where the common unit mas is a milli-arc-sec, or closely 1 part per million of the<br />
full radius).<br />
More rapidly rotating stars can have substantial bulge, which leads to the interesting effect known as the [http://en.wikipedia.org/wiki/Von_Zeipel_theorem von Zeipel theorem] which implies that there is also a substantial dimming of the surface brightness at the equator of such a star.<br />
In the case of the Sun both rotational ellipticity and rotational dimming are tiny and therefore a wonderful challenge for observers.<br />
<br />
Many other potential mechanisms could change the shape of the Sun.<br />
The surface is known to [http://www.aip.de/~mku/difrot/difrot.html rotate differentially] in the sense that the equator region has an angular velocity greater than the polar regions.<br />
There must be internal differential rotation as well, although it is increasingly well constrained by [http://soi.stanford.edu/results/heliowhat.html helioseismology], and what we see at the surface may conceal mass internal mass distributions that can affect the shape.<br />
Such effects are related to [http://en.wikipedia.org/wiki/Tests_of_general_relativity tests of relativity theory], from which Einstein famously explained the anomalous perihelion precession of Mercury.<br />
The observation of this precession was an amazing coup of 19th-century astrometry.<br />
Finally there are links to the [http://en.wikipedia.org/wiki/Solar_cycle solar cycle]; for example a north-south [http://www.noao.edu/noao/staff/irenegh/meridional_web/meridional.html meridional circulation] is required to explain the 11-year cycle - a hot research topic nowadays because of the<br />
recent [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_-_don%27t_panic_yet%21 anomalous minimum].<br />
This could well have detectable effects at the surface of the Sun other than the subtle motions of magnetic elements that trace out the flow associated with the cycle.<br />
<br />
== A history of the measurements ==<br />
<br />
[[Image:126fig2.jpg|left|thumb|300px|Figure 2: Early measurements of the solar oblateness, using both photographic plates and [http://en.wikipedia.org/wiki/Heliometer heliometers]. <br />
These early researchers were interested in the relationship between the oblateness and the phase of the solar [http://www.solarcycle24.com/ sunspot cycle], a question remaining open in spite of the great improvement in measurement accuracy since the 19th century.]]<br />
<br />
Early measurements of solar oblateness used the classical techniques of astrometry such as the measurement carefully exposed photographic plates, or visual observations through transit measurements or via [http://en.wikipedia.org/wiki/Heliometer heliometers] specifically designed for this purpose.<br />
The precision of such measurements (see Figure 2), although well below an arc second, is far from what is known now to be necessary, but that did not keep the observers from speculating about the apparent variations (all due to measurement uncertainty) that they detected - the solid line in Figure 2 compares the observations with the sunspot number.<br />
The [http://en.wikipedia.org/wiki/Astrolabe astrolabe] can also be adapted for measurements of the solar diameter, and can be instrumented with modern detectors as well.<br />
The modern era of solar oblateness measurements can be said to have begun with [http://en.wikipedia.org/wiki/Robert_H._Dicke Dicke], who wished to make measurements sufficiently accurate to test rival theories of [http://en.wikipedia.org/wiki/General_relativity general relativity].<br />
Dicke's ground-based telescope contained several innovations designed to make the measurement differential in nature and as immune as possible to [http://www.schoolsobservatory.org.uk/astro/tels/seeing.shtml "seeing"] conditions.<br />
The device largely succeeded in this, especially after it was moved from Princeton, New Jersey, to a better site in California.<br />
It also inspired other instrumental development.<br />
<br />
Systematic observations of solar oblateness now continue from ground-based observatories, and the measurement errors have dropped to the order of one mas (milli arc sec, or about a part per million).<br />
The new era, as mentioned, began with the observations of Dicke & Goldenberg as reported in [2], but these data were quickly superseded by superior observations, also ground-based.<br />
Figure 2 shows how the error bars dropped during this intermediate period of observation, roughly speaking from 1960 to 1990 (see [1] for details of the individual data).<br />
<br />
[[Image:126fig3.jpg|left|thumb|400px|Figure 3: The intermediate period of specialized, modern ground-based observations of solar oblateness.<br />
The first point on this graph was the revolutionary observation of Dicke & Goldenberg, as reported in [2].]]<br />
<br />
[[Image:126fig4.jpg|right|thumb|400px|Figure 4: The modern data, including the single point thus far published from the RHESSI data.<br />
The black lozenges show the continuing ground-based observations from [http://en.wikipedia.org/wiki/Pic_du_Midi_de_Bigorre Pic du Midi] (see also the Pic du Midi [http://ljr.bagn.obs-mip.fr/histo/histo.html home page]).<br />
The green points show the recent solar maximum; please refer to [1] for fuller details.]]<br />
<br />
It has now become possible to do observations from balloon-borne platforms and from actual satellites in space.<br />
Because the measurements require such great precision, general-purpose solar telescopes are not very useful for oblateness measurements in general, although any full-disk image can be analyzed for its astrometric content.<br />
Thus the observations reported from [http://soi.stanford.edu/ SOHO/MDI] and soon, we anticipate, from [http://hmi.stanford.edu/ SDO/HMI], are lineal descendants of the astrometric plate reductions shown in Figure 1.<br />
Nowadays of course the plates have been replaced by [http://nobelprize.org/nobel_prizes/physics/laureates/2009/press.html CCDs].<br />
<br />
The RHESSI observations [3] presently hold the record in terms of precision, with an oblateness uncerainty of only 0.14 mas. <br />
This has been possible because of a fortunate and accidental imitation of one of Dicke's great ideas: to have the telescope rotate rapidly. <br />
By this artifice many systematic errors become more tractable, since the rotation allows the observer to study deviations from a circular shape almost directly. <br />
The lowest-order deviation is the oblateness term, but as mentioned in the Introduction we fully expect that more precise measurements will reveal higher-order shape terms.<br />
Figure 4 shows the single RHESSI data point, but it is hard to spot at first because of its small error bars.<br />
Eventually we hope to have a full time series of comparable data, extending at least from 2002 to the time of writing (2010).<br />
Other notable data from the current epoch include a continuing series of high-precision measurements from<br />
[http://en.wikipedia.org/wiki/Pic_du_Midi_de_Bigorre Pic du Midi], shown as lozenges with error flags in Figure 4.<br />
The green points in Figure 4 show solar activity, but in terms of the [http://www.scholarpedia.org/article/Solar_activity facular index] rather than sunspot number.<br />
<br />
The RHESSI data, as plotted, show an apparent oblateness greater than 10 mas.<br />
As described in [3], this appears to be the result of solar [http://www.scholarpedia.org/article/Solar_activity magnetic activity] in the form of faculae and enhanced network; the elimination of this contaminating signal brought the apparent oblateness down to the level predicted for uniform rotation.<br />
This result needs confirmation, which we hope to have soon both from RHESSI and from other space missions as well as the ground-based observations.<br />
<br />
== Conclusions ==<br />
<br />
The measurement of solar oblateness has a rich history and a some profound implications touching on relativity theory, for example. <br />
RHESSI is making its contributions to this history [3].<br />
The shape of the Sun is one of the ways we have now for peering into its interior and learning empirically about flows and motions there that would otherwise only be guessed from theoretical considerations.<br />
The paramount of these techniques is of course helioseismology, a technique now being extended to the stars.<br />
<br />
== References ==<br />
<br />
[1] [http://www.sciencedirect.com/science?_ob=ArticleURL&_udi=B6VHB-4YJ4NKR-4&_user=121723&_coverDate=03%2F06%2F2010&_alid=1307462546&_rdoc=2&_fmt=high&_orig=search&_cdi=6062&_sort=r&_docanchor=&view=c&_ct=5&_acct=C000009999&_version=1&_urlVersion=0&_userid=121723&md5=0f904d940638a417b1641698fd30b82f A brief history of the solar oblateness. A review]<br />
<br />
[2] [http://adsabs.harvard.edu/abs/1967PhRvL..18..313D Solar oblateness and general relativity]<br />
<br />
[3] [http://adsabs.harvard.edu/abs/2008Sci...322..560F A large excess in apparent solar oblateness due to surface magnetism]<br />
<br />
[[Has article subject:: history| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Relative_and_(maybe)_Absolute_RHESSI_Detector_Efficiency:_2002-2008Relative and (maybe) Absolute RHESSI Detector Efficiency: 2002-20082018-09-18T02:35:18Z<p>Schriste: /* References */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Nugget Details<br />
|number = 117<br />
|first_author = Jim McTiernan<br />
|publish_date = 22 December 2009<br />
|next_nugget = [[Cycle 24 has begun]]<br />
|previous_nugget = [[A tiny white-light flare]]<br />
}}<br />
<br />
== Introduction ==<br />
<br />
[http://hesperia.gsfc.nasa.gov/hessi/ RHESSI] has now been observing for nearly eight years. How have the [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 detector] sensitivities changed since it was launched? To characterize this important detector property we use a method and software devised by Brian Dennis and Kim Tolbert. This is based on fitting a simple thermal spectrum, for which there are two free parameters, to a given flare observation near the peak time. The parameters are temperature, T, and emission measure, [http://en.wikipedia.org/wiki/Spectroscopy EM]. We do this first for detector 1. Next, thermal components are fit to the spectra for the other eight detectors, with T fixed at the value found for detector 1, T<sub>1</sub>. This gives us a measure of the relative sensitivity of the detectors.<br />
<br />
== The Flare Sample ==<br />
<br />
We want to do this calculation for as many flares as possible. Here we chose every flare on the [http://http://hesperia.gsfc.nasa.gov/hessidata/dbase/hessi_flare_list.txt RHESSI Flare List] which was observed with no attenuators, no particle events, no data gaps, with the spacecraft at low [http://en.wikipedia.org/wiki/Geomagnetic_latitude geomagnetic latitude], and at least 5 minutes from the [http://en.wikipedia.org/wiki/South_Atlantic_Anomaly SAA]. <br />
These are all technical matters details that might bias the results of the cross-calibration and thus make it less precise.<br />
In addition to the flare interval being "clean" in this manner, we also require that the 12 second time intervals before the flare start time and after the flare end time be "clean", so that we can use those times to calculate the background level. This process resulted in a sample of 10,661 flares. Spectra were fit for each flare for the same time range that was used to find the [http://sprg.ssl.berkeley.edu/~jimm/hessi/flare_list_20061004 flare position]. This time interval is usually an interval of 2 minutes or less (depending on flare duration) at the flare peak in the 6 to 12 keV range. The spectra were assumed to be isothermal, and fit for the energy range from 6 to 20 keV, using an energy resolution of 1/3 keV. <br />
<br />
Not all of the flares were fit successfully for all detectors; flares which did not have count rates that were more than 3 sigma above the background level in more than 2 channels were discarded; flares for which the calculated EM was less than 10<sup>43</sup> cm<sup>-3</sup> were discarded; flares with background-subtracted count rates totaled over the 6 to 20 keV range which were less than zero were discarded. These tests were applied for each of the detectors 3, 4, 5, 6, 8, and 9. Spectra from detectors 2 and 7 have been ignored, since those detectors have reduced energy resolution in the 6 to 20 keV range.<br />
<br />
The full sample contains 7,740 flares, from the start of the mission until 1 November 2009.<br />
<br />
== Results for relative sensitivity ==<br />
<br />
[[Image:Hessi_dets_test.png|300px|thumb|center|'''Figure 1''':This is the relative detector efficiency for front detectors 3, 4, 5, 6, 8, and 9 ]]<br />
<br />
Figure 1 shows 60-day averages of the ratio between the EM of each of detectors 3, 4, 5, 6, 8, and 9 and detector 1 for the full time range. From the plot we see that relative sensitivity was closely the same for all detectors until 2006. Early in 2006, the relative sensitivity of detector 3 began to drop. Later in 2006 the detector 5 sensitivity rolls over. Detectors 4, 6, 8, and 9 lose sensitivity during 2007. The first data gap on the plot is for the detector [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 anneal] which took place in November 2007. Post-anneal, the detectors recover sensitivity. The second data gap on the plot is just a quiet time. After the second gap, we can see detectors losing sensitivity again.<br />
<br />
Note that the relative sensitivity of all detectors seems to be greater than unity for the early part of the mission. This may be a systematic effect resulting from holding the temperature fixed at the level found using detector 1 for each of the detectors, i.e. that something about detector 1 leads to a biased temperature value. This possibility is being investigated. <br />
<br />
== What about absolute sensitivity? ==<br />
<br />
For each flare we also measured [http://www.swpc.noaa.gov/today.html#xray GOES] flux averaged for the spectral interval. We expect that the GOES detectors are less susceptible to decay then the RHESSI detectors, which suffer from radiation damage. For most of the flares GOES 10 data were used, but there are gaps in the GOES 10 coverage, so there are flares with GOES 11 and 12 data included. (It turns out that even if non-GOES 10 flares are discarded, we come to the same overall conclusion.) Figure 2 shows 60 day averages of the ratio of RHESSI detector 1 emission measure to GOES fluxes for the full sample, which should show the variation of the relative sensitivity of detector 1 to GOES. Since we see that the other detectors lose sensitivity, we expect that detector 1 should lose some sensitivity, but it seems to gain sensitivity relative to the GOES detectors. How is this possible?<br />
<br />
<br />
{| border="0" align="center"<br />
|+<br />
|-<br />
| valign="top"|<br />
[[Image:Hessi_dets_test1.png|300px|thumb|right|'''Figure 2''':60 day averages of the ratio of RHESSI detector 1 EM to GOES fluxes for the full sample.]]<br />
| valign="top"|<br />
[[Image:Hessi_dets_test0.png|300px|thumb|left|'''Figure 3''':GOES Hi channel flux versus RHESSI detector 1 emission measure for the full sample.]]<br />
|-<br />
|}<br />
<br />
It turns out that there are some effects that are not being accounted for. Figure 3 shows the GOES high-energy, short-wavelength channel flux versus RHESSI detector 1 EM for the full sample. We expect correlation between these two quantities, considering that both instruments are observing X-rays, but with different energy ranges; the GOES energy response is weighted towards energies less than 6 keV, while the RHESSI energy response is small at these low energies and increased at higher energy, where GOES is insensitive. Also the correlation depends on the shape of the observed photon spectrum, which here is assumed to be due to an isothermal source. We are fairly sure that<br />
flares are not isothermal (Ref. 1), and this non-isothermality may be responsible for some of the scatter on the plot.<br />
<br />
There is correlation for relatively large flares, but there is quite a bit of scatter, and for small events, there is no correlation at all. There is a lower limit for GOES, and there are line-like structures in the GOES fluxes which show the effects of digitization in the GOES data. We expect that some of the scatter is due to lack of background subtraction for the GOES data. (Background subtraction for GOES will be included in the next iteration of this project.) This will be a smaller effect for larger flares. Also we may hope to reduce scatter by comparing GOES flux to RHESSI count rate, rather than EM since there is one less level of data processing involved. So for the comparison shown in Figure 4 we restrict the number of flares by only taking relatively large flares, with EM > 10<sup>45</sup> cm<sup>-3</sup>, which is the value below which the correlation between flux and EM disappears. This leaves 2,580 flares. Also we will compare GOES flux with RHESSI count rate, instead of EM. <br />
<br />
Once we have made these changes, we get a good correlation, as shown in Figure 4. The red dashed line shows a linear least-squares fit to the data with a slope of 0.55. There is still scatter; we expect that this will decrease when the GOES background subtraction is included. If we plot the 60-day averages of the ratio of RHESSI detector 1 counts to GOES flux for this flare sample, we obtain the results shown in Figure 5. Note that the error bars in both Figures 2 and 5 are calculated using the standard deviation of the mean of the ratios in each time bin, which is the standard deviation divided by the square root of the number of flares per time bin. There are more flares per bin during the early part of the mission (typically 100 flares/bin) than later (often less than 10 flares/bin); thus the error bars are smaller earlier in the mission. Also there are now three data gaps, one for annealing, and two due to lack of activity. <br />
<br />
{| border="0" align="center"<br />
|+<br />
|-<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct04.png|300px|thumb|left|'''Figure 4''':GOES Hi channel flux versus RHESSI detector 1 6 to 20 keV count rate for the flares with EM greater than 10<sup>45</sup> cm<sup>-3</sup>.]]<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct14.png|300px|thumb|right|'''Figure 5''':60 day averages of the ratio of RHESSI 6 to 20 keV counts to GOES fluxes for the small sample of flares with EM greater than 10<sup>45</sup> cm<sup>-3</sup>.]]<br />
|-<br />
|}<br />
<br />
In Figure 5 the variation of RHESSI counts with respect to GOES flux is much less than the variation shown in Figure 1. The level now looks relatively constant. There is no evidence from this plot that RHESSI detector 1 has lost or gained sensitivity with respect to GOES. It is also probable here that subtracting the GOES background for each flare will affect these results. This will be addressed soon.<br />
<br />
== How does the sample change affect the original plot? ==<br />
<br />
[[Image:Hessi_dets_test_ct4.png|300px|thumb|center|'''Figure 6''':This is the relative detector efficiency for front detectors 3, 4, 5, 6, 8, and 9, but using count rates and only the small sample of 2,580 flares with EM greater than 10<sup>45</sup> cm<sup>-3</sup>.]]<br />
<br />
Using only relatively large flares in the sample, and using count rate rather than EM for the comparison has a small effect on the relative sensitivity plot, as shown in Figure 6. Except for the fact the the detector 3 curve is less than one at the start of the mission, and the odd large value for detector 8 near the end of 2008, not much has changed.<br />
<br />
== Conclusions ==<br />
<br />
Detectors 3, 5, 6, 8, and 9 lose sensitivity relative to detector 1, especially in 2007. Some of that sensitivity is regained by annealing, but it is currently decreasing again.<br />
We will be able to anneal the RHESSI detectors again, and expect as before to restore at least part of their original performance, but we want to do this at the optimum time. <br />
This calibration exercise will help us to understand when that time will come. <br />
<br />
There is no obvious loss of sensitivity in detector 1, relative to GOES. It seems odd that the other detectors can lose sensitivity, so why not detector 1 as well? The next step for this work is to try to subtract pre-flare background for the GOES flare results. Watch this space for updates.<br />
<br />
== Update: add GOES background subtraction <font color=#0000f0> (3 January 2010) </font> ==<br />
<br />
Here are what the results look like when you do subtract GOES background. As a first step, we subtracted pre-flare background values from the GOES fluxes before comparing to RHESSI count rates. For each flare the pre-flare background time interval used is the same as the pre-flare interval found for the RHESSI spectra. (Note that this is slightly different than the background subtraction process for RHESSI, which included both pre- and post-flare background levels and interpolated. Since GOES flares tend to last longer than RHESSI flares, we are concerned that using both the pre- and post flare intervals calculated from RHESSI data may give inaccurate results.) It turns out that only 2055 flares had excess background-subtracted GOES flux. The results for this sample are shown in Figures 7 and 8. In Figure 7, the scatter in the plot of GOES flux versus RHESSI counts is much less than that in Figure 4, as expected. The red dashed line is a least-squares fit to the data, with a slope of 0.86. <br />
<br />
{| border="0" align="center"<br />
|+<br />
|-<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct0b.png|300px|thumb|left|'''Figure 7''':Background-subtracted GOES Hi channel flux versus RHESSI detector 1 6 to 20 keV count rate.]]<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct1b.png|300px|thumb|right|'''Figure 8''':60 day averages of the ratio of RHESSI 6 to 20 keV counts to background-subtracted GOES fluxes.]]<br />
|-<br />
|}<br />
<br />
In Figure 8, which shows the variation of the ratio of RHESSI counts to background-subtracted GOES flux, the relative sensitivity finally shows a decrease. Relative to GOES Hi, RHESSI detector 1 loses about half of its sensitivity by the annealing in November 2007. This is shown a little more clearly in Figure 9, which normalizes the ratio to be 1 at the start of the mission. As in the other figures, sensitivity increases after annealing (first datagap), has been decreasing since, and is now at a level of approximately half of the original sensitivity. <br />
<br />
[[Image:Hessi_dets_test_hi_ct1b.png|300px|thumb|center|'''Figure 9''':This is the relative detector 1 efficiency normalized to be 1.0 for the first 60 days of the mission, using background-subtracted GOES fluxes for the comparison.]]<br />
<br />
== New Conclusion ==<br />
<br />
There is now an obvious loss of sensitivity in detector 1, relative to GOES, which we find by including pre-flare background subtraction for the GOES flare results.<br />
<br />
== References ==<br />
1. http://adsabs.harvard.edu/abs/1999ApJ...514..472M The Solar Flare Soft X-Ray Differential Emission Measure and the Neupert Effect at Different Temperatures by J.McTiernan, G. Fisher and P. Li<br />
<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has article subject:: RHESSI| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Imaging_through_visibility_interpolation:_uv-smoothImaging through visibility interpolation: uv-smooth2018-09-18T02:34:47Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Imaging through visibility interpolation: uv-smooth<br />
|title = Nugget Details<br />
|number = 113<br />
|first_author = Anna Massone and<br />
|second_author = Michele Piana<br />
|publish_date = October 26, 2009<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::114]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index:112]]}}<br />
}}== Introduction ==<br />
<br />
RHESSI imaging works by [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=8 image modulation],<br />
rather than by focusing.<br />
The raw data thus do not reveal the image immediately, but only after processing.<br />
This procedure is mathematically (but not physically) analogous to the manner in which an<br />
[http://en.wikipedia.org/wiki/Interferometry interferometic]<br />
[http://en.wikipedia.org/wiki/Radio_astronomy radio telescope]<br />
images the sky.<br />
<br />
Straightforward, ''traditional'' ways of making images from RHESSI data work much as you would use PhotoShop to enhance the image from a digital camera: you start with a [http://scienceworld.wolfram.com/physics/DirtyMap.html "dirty map"] of the flare and use deconvolution software (e.g., <br />
[http://www.britannica.com/EBchecked/topic/120779/CLEAN CLEAN], <br />
[http://www.adass.org/adass/proceedings/adass98/puetterrc/ PIXON])<br />
to reduce the blurring effects of the collimators and so reconstruct a higher quality, less noisy, image of the real scene. This Nugget describes a recently-developed alternative approach, which is a little bit more sophisticated but which is tailored to the specific way in which RHESSI encodes imaging information. The rotation of the grid pairs of each of RHESSI’s <br />
[http://en.wikipedia.org/wiki/Rotational_Modulation_Collimator Rotating Modulation Collimators] (RMCs) provides a temporal modulation of the arriving flux; the pattern of this temporal modulation provides not so much an image but rather a specific set of spatial Fourier components of the source. These data, the purest form of the RHESSI data stream, are called visibilities – they are calibrated measurements of specific spatial <br />
[http://en.wikipedia.org/wiki/Fourier_transform Fourier components] of the emitted radiation field. Visibilities are complex numbers with quantifiable (complex) error bars.<br />
<br />
[[Image:Fig1_uv_smooth.jpg|350px|thumb|right|'''Figure 1''': Top left panel: the visibilities (magnitudes). Top right panel: blurred map obtained by inverse Fourier transforming the visibilities. Bottom left panel: visibility surface (magnitude) from interpolation. Bottom right panel: uv<sub>-</sub>smooth reconstruction.]]<br />
<br />
With its nine <br />
[http://en.wikipedia.org/wiki/Rotational_Modulation_Collimator RMCs], RHESSI produces visibilities on a set of nine concentric circles in the spatial frequency plane, each circle corresponding to the spatial frequencies sampled by a particular RMC as the spacecraft rotates. Because the grid pitches of RHESSI’s RMCs are arranged in a geometric progression with factor <math>\sqrt{3}</math>, the radii of these circles in the spatial frequency plane are also distributed in a geometric progression with the same factor; they range from 0.0027 arcsec<sup>-1</sup> for the coarsest grids to 0.221 arcsec<sup>-1</sup> for the finest grids. Armed with these visibilities, one can still reconstruct meaningful images by simply inverse Fourier transforming them from the frequency to the physical plane. However, such an approach, while at a first sight straightforward, is severely limited by RHESSI’s relatively sparse sampling of the frequency plane – there are simply not enough Fourier components available to make a high quality image.<br />
<br />
<br />
Enter ''uv<sub>-</sub>smooth''. This new imaging technique uses a couple of simple but powerful mathematical tools to circumvent the problems of RHESSI’s sparse frequency-space sampling. The first of these is '''interpolation''': uv<sub>-</sub>smooth generates a smooth continuum of visibilities within the disk encompassing the sampled area of the spatial frequency plane. The interpolation technique is a ''two-dimensional spline fit''; it finds the ''locally minimally bent'' smooth surface passing through all experimentally sampled frequencies. (Think of a set of vertical tent poles on a set of nine circles in the ground [the measured visibilities]; uv<sub>-</sub>smooth finds the shape of the tent canvas that fits smoothly over these poles.) Once we have this continuous visibility surface, we can take the interpolated visibility values at regularly spaced points in the Cartesian spatial frequency plane and apply a very efficient <br />
[http://en.wikipedia.org/wiki/Fast_fourier_transform Fast Fourier Transform] (FFT) technique to produce an image. Since this image is constructed from a relatively large number of Fourier components, it will be of considerably higher quality than the image produced using only the directly sampled visibility values.<br />
<br />
The second crucial step in the method uses an iterative scheme that incorporates a '''positivity''' constraint - at each step the negative intensity values in the image (that are associated with ''ringing'' of each of the Fourier components) are removed and the Fourier transform of this modified image is taken. This new set of Fourier components is then smoothed as before and the corresponding image generated. As a result of this iteration scheme, additional Fourier components outside the sampled disk are generated. This effectively extrapolates the sampled RHESSI data set to higher spatial frequencies (i.e., smaller angular resolutions), so endowing the method with the property of super-resolution.<br />
<br />
== An illustration of the method ==<br />
<br />
In the top left panel of Figure 1, the spikes (''tent poles'') you can see are the absolute values (magnitudes) of the (complex) visibilities that would be directly measured if RHESSI were to observe a simple two-dimensional Gaussian source. Processing these directly measured visibility values using an inverse Fourier transform results in a rather dirty, blurred reconstruction of the original Gaussian source (top right panel). Using uv<sub>-</sub>smooth, we replace the sparse (''tent pole'') data with a smooth continuum (''canvas'') obtained via spline interpolation: the measured visibilities are maintained while additional visibilities are added in between. The result is the surface of visibility magnitudes in the bottom left panel of Figure 1; a similar surface exists for the complex phase. This surface provides information inside the disk in the spatial frequency plane that is spanned by the measured visibilities, but with much more complete coverage. Computing the inverse Fourier transform of this smooth visibility surface using an <br />
[http://en.wikipedia.org/wiki/Fast_fourier_transform FFT] routine results in an image of the source; the iterative scheme discussed above is then applied to reduce the appearance of negative values in the image. The result is shown in the bottom right panel of Figure 1; the recovered image rather faithfully reproduces the form of the original source.<br />
<br />
[[Image:Fig2_uv_smooth.jpg|350px|thumb|left|'''Figure 2''': Left panel: visibility surface (magnitude) from interpolation. Right panel: visibility surface (magnitude) corresponding to the uv_smooth reconstruction: the extrapolation effect is induced by the positivity constraint.]]<br />
<br />
Note that the influence of the positivity constraint in the iterative algorithm that forms the second part of the method not only reduces ringing effects in the image but also has the rather ''magical'' consequence (justified by a beautiful and rigorous mathematical theory) that the visibility surface does not go to zero abruptly, as it does when using only data within the region of the spatial frequency plane sampled by RHESSI (left panel of Figure 2). Instead, as shown in the right panel of Figure 2, it drops to zero via a smooth tapering that effectively provides an extrapolation of the data in the spatial frequency plane and so (consistent with information theory) super-resolution in the image plane. This explains why uv<sub>-</sub>smooth is so accurate in reproducing the features (size, shape) of the original source.<br />
<br />
<br />
== A real flare ==<br />
<br />
[[Image:Fig3_uv_smooth.jpg|500px|thumb|right|'''Figure 3''': 2002 February 20, 11:06:02-11:06:34 UT, event (energy range ε = 22 - 26 keV). Left panel: uv_smooth reconstruction. Middle and right panels: average of 100 uv_smooth reconstructions from randomly perturbed visibilities with the corresponding standard deviation.]]<br />
<br />
It is very easy to apply uv<sub>-</sub>smooth to real visibilities measured by RHESSI. You take the visibilities, interpolate them, apply the iterative algorithm and after a few iterations (the scheme typically converges in less than 20 steps) you have a nice image of the flare. Everything is fast, reliable and robust. A nice example is in Figure 3, which uses data from the [[Has event date::2002 February 20 11:06:02]] to 11:06:34 UT, event (energy range ε = 22 - 26 keV). In the left panel we show the uv<sub>-</sub>smooth reconstruction; in the middle and right panels we exploit the incredible speed of the method to perform a Monte Carlo experiment using 100 random realizations of the input data. The average image (center panel) and its standard deviation (right panel) are obtained in less than 30 seconds of computational time on a standard PC.<br />
<br />
== Conclusions ==<br />
<br />
uv<sub>-</sub>smooth is a very natural way for making images with RHESSI. It uses the RHESSI RMC data in its ''purest'' form, coupled with a fast FFT-based algorithm that makes image processing with uv<sub>-</sub>smooth extremely fast and reliable. Many experiments using simulated data have shown that it is applicable to a wide variety of source structures, reliably reproducing both the photometric intensity and shape of individual compact sources, and the intensity ratios of subsources within complex flaring structures. uv<sub>-</sub>smooth has been recently added to the RHESSI GUI. Try it! You won’t be disappointed.<br />
<br />
== Reference ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2009ApJ...703.2004M Hard X-ray Imaging of Solar Flares using Interpolated Visibilities]<br />
<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has article subject:: education]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Browse_the_RHESSI_data!Browse the RHESSI data!2018-09-18T02:30:17Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Browse the RHESSI data!<br />
|number=20<br />
|first_author=Albert Shih<br />
|publish_date=16 January 2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::21]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::19]]}}<br />
}}<br />
<br />
==Introduction==<br />
<br />
RHESSI data (and software) are of course freely available to all, but to view the data directly often requires some training or study. Accordingly one needs a simple overview, and [http://sprg.ssl.berkeley.edu/~tohban/browser| Browser] - a Web interface - fills this need. This science nugget will not be providing knowledge on some scientific topic, but rather on how to use Browser to its maximum potential in the pursuit of science. As many of you already know, Browser is a very useful tool to browse through many different types of RHESSI-related data plots simultaneously. However, I suspect that many of you may not realize some of its not-so-obvious capabilities. If so, read on!<br />
<br />
==Ground Rules==<br />
<br />
First, if you are rather unfamiliar with Browser, just play with it! Perhaps pick a time when you know something interesting happened, and start clicking around on various icons and links. Feel free to experiment!<br />
<br />
The most useful icons or links may be the question-mark icon at the top and the "Help" link at the bottom. Both will open the help page, which may answer any basic questions you may have. The next most useful icons are the magnifying-glass icons that are littered across the page. Clicking on such an icon will reveal advanced options that may turn out to be indispensable for your method of using Browser. Finally, one should not forget the red-arrow icons, which indicate a particular time on all plots that have a time axis. Combined with the ability to click on plots to select a specific time, one can click on a feature of interest on one plot, and then the red arrow will identify the same time on all the other plots.<br />
<br />
I will now present two case studies for how one might use Browser. <br />
<br />
==Case Study 1: X4.8 Solar Flare on 2002 July 23==<br />
<br />
[[File:20f4.png|center|thumb|400px]]<br />
<br />
I will first discuss the typical use of Browser to look at solar flares, and I have chosen the much-studied X4.8 solar flare on 2002 July 23. Click on the image on the left to open the actual Browser page in a new window. Now to explain its elements...<br />
Note that there are quite a few plots open, although these are only a fraction of the possible number. Unsurprisingly, the more plots that are open, the longer the page takes to load completely, so I would suggest a smaller set for casual browsing, and then opening other plots when necessary. However, if you ordinarily view only one plot at a time in Browser, I would strongly encourage you to try browsing with multiple plots open.<br />
<br />
The first plot [1] shows the GOES lightcurves with the RHESSI times for eclipse and radiation-belt passages superposed (no data and bad data, respectively). For many flares, one can use these plots to determine whether a given flare has observed. This plot is also useful to jump between multiple flares on a single day (see the red-arrow icon discussion above).<br />
<br />
Next, the [http://hesperia.gsfc.nasa.gov/rhessidatacenter/quicklook/data.html| quicklook] lightcurve plots [2] show the corrected and uncorrected rates for the orbit. "Quicklook" is rich collection of reduced secondary databases - lightcurves, images, and spectra. A brief digression: when you see interesting behavior on the corrected-rates plot, be sure to check the uncorrected-rates plots to make sure that the behavior was not simply a correction artifact. For example, the attenuation by the shutters was corrected for in this flare, but note that the 7000-20000 keV band has picked up a small spike at 0041 UT that is clearly a correction artifact (there is nothing to correct for at such high energies).<br />
<br />
The next quicklook plot [3] shows a sampling of the data from the particle counters on RHESSI (on board mainly to track [http://www-spof.gsfc.nasa.gov/Education/Iradbelt.html| Earth's radiation belts]). The top and bottom halves show lightcurves for the detectors that usually respond to electrons and protons, respectively. In this particular example, there is a short burst of activity after 0100 UT. By using the red-arrow icons to identify the same time on the quicklook lightcurve plots, we see the particle activity did in fact produce an effect visible in the RHESSI lightcurves, although fortunately well after virtually all the interesting flare emission.<br />
<br />
The next three plots [4] are various types of quicklook images for this flare, and the following plot [5] (gaudy and complicated) combines GOES, RHESSI, and [http://www-spof.gsfc.nasa.gov/istp/wind/| WIND] particle data, and also includes the RHESSI data in spectrogram form.<br />
<br />
The final three plots [6] above are examples of monitor-rates plots. The first plot shows the "slow valid" (i.e. good) counts in each of the nine front segments. Of course, the quicklook lightcurve plots are usually much more useful, but this plot has its uses. For example, the artifact in detector 8 due to interference from the aft antenna can be clearly seen (e.g. 1650 UT). The following two plots show the livetimes in the front segments and the rear segments. Note that since flare photons tend to stop in the front segments before reaching the rear segments, the rear segments have substantially more livetime than the front segments. However, this relationship does not hold for non-solar photons such as those from nearby particle precipitation events. There are many more monitor-rate plots, but they are probably too technical for most users.<br />
<br />
==Case Study 2: Magnetar Burst on 2004 December 27==<br />
[[File:20f5.png|center|thumb|400px]]<br />
<br />
Now let us take a look at an event that was not a solar flare, specifically the [[Neutron starquake shakes RHESSI|magnetar burst]] on 2004 December 27. (Even though RHESSI flagged this event a solar flare, it is most definitely not one.)<br />
Looking at the quicklook plots [1], one can see an egregious correction artifact where many of the lightcurves shift to a new baseline after the magnetar burst. Again, one should always remember to check the uncorrected-rates plots when seeing strange behavior.<br />
<br />
The next two quicklook plots [2] show the count lightcurves for the nine front segments and nine rear segments, respectively. One can use these plots in much the same way as the corresponding monitor-rates plots, and, in addition, flares are marked by lines on these plots.<br />
<br />
The next plot [3] is an attempt to display a quicklook image plot for this event. Since this non-solar event cannot be imaged, an error message appears instead.<br />
<br />
Finally, there is the GOES/RHESSI/WIND plot [4]. Note that the GOES data shows that the burst went virtually unrecorded in the low-energy channel. It would have been difficult to pick out this event on the usual GOES-with-RHESSI-times plots. Also note the appearance of the burst in the spectrogram, and compare it against the totally unrelated particle precipitation event that occurred an hour later at 2340 UT.<br />
<br />
==Advanced Features==<br />
<br />
The first advanced feature that you may not have been aware of is that Browser automatically synchronizes links to corresponding data on other websites at the same time as currently selected in Browser. For example, if you have the first case study open, you can click on the [http://www.solarmonitor.org/index.php| Active Region Monitor] link on that page to bring up SOHO data for the 2002 July 23 flare.<br />
Another advanced feature is the ability to create links to a specific time with particular plots open (like the two case studies above). Once you have arranged Browser in the way you want, you can click the "Direct URL" button (or right-click and save the link's location) to access the custom shortcut. One can then bookmark the page, or perhaps send it to a colleague for consultation. Incidentally, if one clicks the "Reset Date" button immediately before using the "Direct URL" button, then the time information is erased, and the direct URL stores only the current view layout (e.g. my preferred view).<br />
<br />
Finally, there is a major timesaver in the form of text fields for the flare number or time (click the magnifying-glass icon next to the year-selection box). Although these fields are somewhat useful as output fields, their real power comes from using them as input fields. One can enter a specific flare number or a specific time (in most [http://www.lmsal.com/solarsoft/| Solarsoft-compatible] formats), and Browser will jump to the corresponding time. This feature works great when analyzing a list of events. Note that there is an eraser button provided to aid copy-and-pasting on systems where selecting automatically overwrites the contents of the clipboard. <br />
<br />
==Concluding Remarks==<br />
<br />
I hope this discussion introduced you to Browser and some new facets that you had not been aware of. Provide me with comments and suggestions on Browser, and if I happen across any free time, I will see what I can do. But, as is, Browser has helped my own research immensely, and I hope that you will be able to say the same.<br />
<br />
'''Biographical note''': Albert Shih is a graduate student at UC Berkeley, working on a PhD thesis analyzing solar gamma-ray flares.<br />
<br />
[[Has article subject:: education]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/X-ray_and_gamma-ray_spectroscopyX-ray and gamma-ray spectroscopy2018-09-18T02:29:56Z<p>Schriste: /* Remarks */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=X-ray and gamma-ray spectroscopy<br />
|number=27<br />
|first_author=Hugh Hudson<br />
|second_author=David Smith<br />
|publish_date=22 January 2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::28]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::26]]}}<br />
}}<br />
<br />
==Introducing one of RHESSI's main purposes==<br />
<br />
Reuven Ramaty, whose name adorns our satellite, specialized in astronomical gamma-ray spectroscopy. A reader may wonder how a photon so energetic as to penetrate inches of lead shielding can ever be tamed spectroscopically. RHESSI does this better than any previous astronomical instrument, and this nugget introduces the subject. RHESSI both images and makes high-resolution spectra simultaneously ("imaging spectroscopy"), but most of our previous nuggets have emphasized the imaging. We hope to have several future nuggets on some of the many observational results of the RHESSI spectroscopic capability.<br />
==Germanium detectors==<br />
<br />
The key to RHESSI's spectroscopic success lies in the properties of "hyperpure" germanium, which has the property of converting energy deposition precisely and linearly into electric charge, which can then be amplified and processed. Such a counter has a prototype in the "proportional counter", which introduced to physics the idea of "non-dispersive spectroscopy," in which one measures the energy of individual photons rather than using prisms or gratings to sort the energies. Non-dispersive spectroscopy, although usually of lower resolution, has the multiplex advantage, meaning that RHESSI can measure all photons of any energy at all times; no spectral scanning is required.<br />
<br />
So how do germanium (Ge) solid state detectors work? RHESSI carries nine large hyperpure Ge crystal detectors. When a photon deposits its energy within the detector, electron-hole pairs (charges) are produced. To sweep up these charges, one needs a strong electric field. The figure below shows the pattern of electric field lines resulting from the high voltage (-4000 V) applied to the outer surface of the detector:<br />
<br />
Figure 1. The electric field in a RHESSI germanium (Ge) detector. The shape of the detector is that of a hollowed-out cylinder. The high voltage is applied at the outer (P+) contact, while the inner (N+) contact is split into two parts to collect electrons from the upper and lower regions of the crystal. The top of the detector points to the Sun. <br />
The electric field configuration effectively splits the detector into two separate sections, the front and rear. An interruption in the inner contact allows these to be read out through separate chains of electronics. One RHESSI Ge detector acts like two! The front detectors catch the lower energy (not so penetrating) X-rays, while very energetic (penetrating) gamma rays travel right through the front and are caught by the rear section. The front-back geometry of the segmentation thus serves to shield the rear part of the Ge from the flood of lower-energy X-rays always produced by a solar flare.<br />
<br />
==Photon Energy==<br />
<br />
The main reason to use germanium -- an expensive detector material that must be operated below 100K -- is its superb energy resolution. Photon energy can be measured to within about 1 keV for hard X-rays in the front segments (below about 200 keV) and to within about 2 to 5 keV in the rear segments. For an MeV photon, this is a very high accuracy! It not only allows solar gamma-ray lines to be identified with specific nuclear processes, it even allows their Doppler profiles to be studied, in analogy to optical spectroscopy. Figure 2 shows an intrinsically narrow solar gamma-ray line from RHESSI's first gamma-ray line flare.<br />
<br />
<br />
Figure 2. Solar line from neutron capture during deuterium formation, from the 23 July 2002 GOES X4.8 flare observed by RHESSI. The natural line is extremely narrow; the width observed is the instrumental resolution of RHESSI's germanium detectors.<br />
When a very energetic photon interacts with the Ge, several things can happen. The simplest process is that all of the energy of the photon is deposited inside the germanium volume (via the mechanisms of photoelectric absorption, Compton scattering, and electron/positron pair production). If one were to plot detected count energy versus true photon energy then this response would be diagonal. Unfortunately, life is not always that simple; often, you don't get all the energy of the photon. The following two images show the response matrix for the front and rear detectors.<br />
<br />
<br />
Figure 3a. A visual representation of the response matrix for the front detectors (left) and rear detectors (right). Note the energy range difference in the axes, since the fronts detect low energies (X-rays) while the rears detect the high energies (gamma-rays). The color display is logarithmic. Click for an image for a larger version.<br />
The non-diagonal response is quite complex and includes many different processes, though most of the response at low energies is actually in the "diagonal elements" of the matrix, the thin yellow line. Above a few hundred keV, the off-diagonal response makes up the majority of the counts.<br />
<br />
Here are these plots again with specific features marked.<br />
<br />
<br />
Figure 3b. Same as figure 3a but with non-diagonal features marked. Click an image for a larger version.<br />
<br />
In cases where the full energy is not collected, the incoming photon does one of the following:<br />
<br />
<br />
#Compton-scatters off of an electron and exits the detector leaving only a fraction of its energy. This accounts for most of the off-diagonal activity and the strong Compton edge.<br />
#Compton-scatters through a small angle in the grids (or anything else) and into the detector.<br />
#Compton-scatters through a large angle (backscatters) into the detector.<br />
#excites K-shell fluorescence peaks from nearby passive material, including the tungsten collimator grids (59 keV, 67 keV).<br />
#generates a germanium K-shell (~ 10 keV) fluorescence photon which escapes.<br />
#generates an electron and positron pair outside of the detectors; the positron subsequently annihilates on another electron, generating the line at 511 keV.<br />
#generates an electron and positron pair inside of the detector, with one or two of these annihilation photons escaping, giving two spectroscopic escape peaks.<br />
<br />
==Remarks==<br />
<br />
The purpose of this RHESSI science nugget was just to introduce some of the ideas behind X-ray and gamma-ray spectroscopy, which is a somewhat obscure subject in astrophysics because of its extraordinary difficulty. In spite of this difficulty, the off-diagonal effects can be rather well accounted for by a combination of simulations and laboratory measurements. Removing the effects of the off-diagonal response is part of the RHESSI software. Future science nuggets will discuss some of the observations made with this system.<br />
<br />
[[Category:Nugget needs figures]][[Category:Nugget need cleaning]]<br />
<br />
Biographical note: Hugh Hudson is at UC Berkeley, and David Smith is at UC Santa Cruz. We thank Richard Schwartz and Albert Shih for the lovely images of the response matrices.<br />
<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has article subject:: RHESSI]]<br />
[[Has article subject:: education]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Radiation_damageRadiation damage2018-09-18T02:29:19Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Radiation damage<br />
|number=34<br />
|first_author=Steven Christe<br />
|second_author=Albert Shih<br />
|publish_date=28 January 2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::35]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::33]]}}<br />
}}<br />
<br />
==Introduction==<br />
<br />
At RHESSI's heart are its nine hyperpure germanium detectors. Space is not only a harsh environment for living organisms but also for radiation sensors (for mostly the same reasons). High-energy charged particles zip along at speeds close to the speed of light and can deposit a lot of energy into materials, causing radiation damage as described below. These particles can come from a variety of sources: from the Sun in the form of solar energetic particles (SEPs...as discussed in a ), from the radiation belts, or from the all-pervading cosmic rays. As previously mentioned RHESSI is only lightly shielded, which makes it more susceptible to radiation damage. RHESSI's orbit takes it through the South Atlantic Anomaly where particle fluxes are very high, not to mention particle events which frequently hit RHESSI. All of these energetic particles can interact with RHESSI's germanium detectors and damage them.<br />
<br />
==Radiation Damage==<br />
<br />
To refresh everyone's memory, energy deposited (say, by a photon) in a semiconductor creates electron-hole pairs. These charges are then collected at opposite contacts through the use of a large voltage applied across the active volume. High energy radiation can create large disordered regions (on the order of 100 angstroms in size) in the germanium crystal, and these regions tend to accumulate negative charge. As a result, these regions act as hole traps and prevent holes from reaching the cathode. If a hole is trapped, then only a fraction of the energy of the incident photon is collected. This leads to tailing and a decrease in energy resolution.<br />
<br />
Another effect of radiation damage (if allowed to go on for a while) is a loss of sensitivity. As the detector becomes more damaged, it becomes difficult to fully deplete the detector, and thus there is a reduction of effective area as there is less active volume. Both of these effects can be seen in RHESSI (see below)<br />
<br />
<br />
Figure 1. A depiction of the radiation-damaged lineshape in one of the RHESSI detectors for observing a monoenergetic 1-MeV photon in the years since RHESSI's launch.<br />
<br />
Figure 2. A comparison between the detection of the background line at 1370 keV between launch and four years afterwards, in the rear segments of each of the nine detectors. Note the line broadening due to radiation damage, and the lack of appreciable active volume in many detectors. As an aside, detector 2 did not have a function rear segment at launch.<br />
==Annealing==<br />
<br />
Fortunately, RHESSI is not lost! One can repair radiation damage through the procedure known as annealing. The detectors are heated (to approximately 100 degrees celsius) for a certain period of time (at least a week) and the damage can be fixed! This procedure increases the mobility of the germanium atoms sufficiently to break up the previously mentioned large disordered regions into smaller clusters that no longer trap holes. Unfortunately, there is additional concern with annealing the RHESSI detectors compared to typical germanium detectors. RHESSI's detectors are each split into two segments, and the contacts for these segments are created by doping the inner surface with lithium. The contacts are normally separated by a region without any doping, but annealing the detector can cause the lithium to diffuse to the point where the contacts touch. In that event, the two segments will fuse into one.<br />
<br />
Test annealing is currently being done down here at SSL on a spare RHESSI detector to determine the how long we can anneal the detectors before they desegment. If the RHESSI detectors become single segments, then their sensitivity to gamma-rays (from solar or cosmic origins) becomes reduced, although probably still better than their present sensitivity. We hope to be able to anneal the RHESSI detectors soon and make them new and fresh again!<br />
<br />
'''Biographical note''': Both Steven Christe and Albert Shih are currently graduate students at U.C. Berkeley.<br />
<br />
[[Has article subject:: RHESSI]]<br />
[[Has observation by:: RHESSI| ]]<br />
<br />
<br />
[[Category:Nugget needs figures]][[Category:Nugget need cleaning]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_VisibilitiesRHESSI Visibilities2018-09-18T02:28:40Z<p>Schriste: /* Ackowledgements */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = RHESSI Visibilities<br />
|title = Nugget Details<br />
|number = 39<br />
|first_author = Ed Schmahl<br />
|second_author = Gordon Hurford<br />
|publish_date = 11-Jun-2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::40]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::38]]}}<br />
}}<br />
<br />
== Visibilities à la Fourier ==<br />
<br />
RHESSI is a Fourier imager, which means that it obtains amplitudes and phases (a.k.a. "visibilities") of the X-ray image modulations. This particular choice for an imager [http://hesperia.gsfc.nasa.gov/~kim/Science.html#imaging_system design] was driven by many forces, not the least of which was the mantra "[http://www.astrosociety.org/pubs/mercury/9504/tenerelli.html faster, better, cheaper]" which put strong cost and size constraints on the [http://sunland.gsfc.nasa.gov/smex/ Small Explorer (SMEX)] satellites of the 1990s. Previous solar instruments that used Fourier techniques were [http://www.daviddarling.info/encyclopedia/H/Hinotori.html Hinotori], [http://www.astro.umd.edu/~ed/gifs/heidi_payload.gif HEIDI], and [http://gedas22.stelab.nagoya-u.ac.jp/HXT/ HXT], but RHESSI has unique new Fourier/visibility-based [http://hesperia.gsfc.nasa.gov/rhessidatacenter/imaging/imaging_concepts.pdf capabilities] that are only now beginning to be exploited.<br />
<br />
What does Fourier have to do with imaging? The 18th century mathematician [http://en.wikipedia.org/wiki/Joseph_Fourier Jean Baptiste Fourier] discovered the principle that any signal can be broken up into a [http://mathworld.wolfram.com/FourierSeries.html superposition of sines and cosines], or the equivalent amplitudes and phases, but he did not himself invent Fourier imaging. Applications of Fourier methods to imaging had to wait until the late 19th century when the amplitude and phase pairs of optical interferometry (see this [http://www.astro.lsa.umich.edu/~monnier/Publications/ROP2003_final.pdf 69-page .pdf document] for full detail) came into use, and these paired variables came to be known as visibilities, as in "fringe visibility". Later in the 20th century [http://www.geocities.com/capecanaveral/2309/page3.html radio] interferometry used the same entities, and now hard X-ray astronomers have joined the visibility game (see this [http://sprg.ssl.berkeley.edu/~ghurford/VisibilityGuide.pdf 8-page .pdf] document for details).<br />
<br />
== How do we create visibilities? ==<br />
<br />
To get actual sine and cosine components of hard X-ray images is difficult, but RHESSI uses a good approximation: triangular waveforms. In the ideal case - constant background, thin grids, and high photon rates - the modulated waveform generated by photons incident on one of RHESSI's sub-collimators has a triangular profile (Figure 1). The amplitude of the waveform is proportional to the intensity of the beam, and its phase and frequency depend on the direction of incidence. For complex sources, and over small rotation angles, the amplitude and phase of the waveform provide a direct measurement of a single Fourier component of the angular distribution of the source. Different Fourier components are measured at different rotation angles and with grids of different pitches.<br />
<br />
[[Image:39_triangle_waveform.png|600px|thumb|center|'''Figure 1''': Ideal hard X-ray grid response (solid curve), with its associated fundamental sinusoid (dashed curve).]]<br />
<br />
== From time tags to visibilities ==<br />
<br />
We illustrate the end-to-end process of computing visibilities for a particular flare interval, a single subcollimator, and a single energy band. The primary RHESSI data are photon time tags and pulse heights, which must be transformed in various ways before imaging is possible. RHESSI's time tags and pulse heights for a selected 100-ms interval are shown in Figure 2a. These time tags have not yet been binned in energy or time, so they are difficult to interpret by eye. Several gaps produced by cosmic-ray hits are present. <br />
<br />
[[Image:39_time-tagged_pulse_hts.png|600px|thumb|center|'''Figure 2a''': Time-tagged pulse heights]]<br />
<br />
A sequence of basic (linear) operations must be performed on these data to convert them to visibilities. The next step is to bin the data in energy and in uniform roll-angle bins, as shown in Figure 2b.<br />
<br />
[[Image:39_energy-binned_pulse_hts.png|600px|thumb|center|'''Figure 2b''': Energy-binned pulse heights]]<br />
<br />
The next step is to "histogram" the time tags into "roll bins," ie units of RHESSI rotational phase, as shown in Figure 2c for 3 rotations.<br />
<br />
[[Image:39_modulated counts.png|600px|thumb|center|'''Figure 2c''': RHESSI modulation profile. The green vertical lines mark the boundaries of spacecraft rotations, (units of 360 degrees) and the red tick marks indicate the roll bins selected for phase-bin stacking. Three sets of rollbins are identified by the colors blue, orange, and cyan.]]<br />
<br />
We now progress further towards constructing visibilities by binning the modulation cycles into roll bins, which are then binned again into 12 aspect phase bins. In this case there are 16 roll bins, three of which are outlined in color (blue, orange, and cyan) in each rotation. After phase binning, each of these rollbins can be co-added.<br />
<br />
[[Image:39_stacked_counts_3_rbins.png|600px|thumb|center|'''Figure 2d''': Stacked roll bins 3,4 & 5 taken from the three selected regions shown in corresponding colors in the previous figure.]]<br />
<br />
This process, called "[http://hesperia.gsfc.nasa.gov/~schmahl/the_stacker/xdoc_hsi_phz_stacker.png stacking]" increases the signal-to-noise ratio and provides a platform for computing a mean amplitude and phase. This generally improves the aspect phase coverage because the angular drift of the telescope causes the aspect phase to differ from one rotation to the next. After stacking we fit sinusoids to the profiles (red curves in Figure 2d).<br />
<br />
In order to calibrate the parameters generated by the fits, a number of corrections must be made. The grids have a finite thickness, which produces internal shadowing and modifies the triangular-like response; this also causes a decline of transmission as a function of the off-axis distance; the grid slats are not all perfectly aligned parallel to the optical axis, which leads to a "venetian blinding" effect, which makes even and odd half-rotations unequal; and there are detector-to-detector sensitivity differences. Fortunately, RHESSI was calibrated in many ways before launch, and the instrument has many self-calibration capabilities, so all of these effects can be calibrated out. Using the RHESSI calibrations, one may remove the instrumental dependence and convert the sinusoids to photon rates.<br />
<br />
The final steps in constructing visibilities are to output the amplitudes and phases of the fitted sinusoids, and compute error estimates based on photon statistics and instrumental systematics. Figure 2e shows the computed amplitudes in the upper panel, and the visibility phases in the lower panel. <br />
<br />
[[Image:39_amplitudes_and_phases.png|600px|thumb|center|'''Figure 2e''': Visibility amplitudes and phases derived from the stacked counts in Figure 2d. Roll bins 3,4 & 5 taken from the three selected regions are shown in the same colors as the previous figure.]]<br />
<br />
Note that the visibility amplitudes are in units of photon flux, unlike the stacked modulation profiles, which are in counts per bin. Since the calibrations have been applied, the visibilities are very nearly independent of the instrument. The importance of this instrumental independence is hard to overestimate. It makes it possible, for example, to use radio imaging techniques such as [http://hesperia.gsfc.nasa.gov/~schmahl/MEM_NJIT/mem_njit_tutorial.html MEM] for RHESSI imaging, and to do forward fitting ([http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=35 see earlier nugget]). We list several other applications of visibilities later.<br />
<br />
== RHESSI visibility amplitudes ==<br />
<br />
Above, we showed data for only a single subcollimator. But RHESSI's great imaging power results from its coverage of a wide range of spatial scales with its nine subcollimators. Here (Figure 3) we show the amplitude dependence as a function of both subcollimator (SC) and roll angle (PA). The X-axis is a linear combination of both variables in the form SC+PA/180. Thus, for example, the roll angles 0 to 180 degrees are shown for subcollimator 6 for X between 6.0 and 7.0. <br />
<br />
[[Image:39_ed.gif|600px|center|thumb|'''Figure 3''': Observed visibility amplitudes (blue crosses) for a flare interval as a function of subcollimator (SC=1-9) and position angle (PA=0-180°) of the grids. Each of the 9 vertical panels shows the amplitude as a function of PA for one subcollimator (labeled by red digits below the X axis). The red curve represents a model using two Gaussian sources, and the green squares show the residuals relative to the model. For a given subcollimator (6 and 7 are good examples), the amplitude rises and falls sinusoidally while the grids rotate from PA=0 to PA=180°. In general, such sinusoidal variation indicates an extended or double source.]]<br />
<br />
The observed amplitudes are indicated by blue crosses and their associated estimated standard errors are shown by vertical error bars. Note the variable amplitude as a function of roll angle for each subcollimator. This variation is caused by the "beating" of two flare sources against each other. Note also the gradual falloff of amplitude as the subcollimators become finer (towards smaller X). This falloff is the result of the finite sizes of the sources. When the angular pitch is smaller than the source size, the modulation amplitude is reduced (and the error bars become larger).<br />
<br />
The solid red curve shows the visibility computed for a model given by two Gaussian sources. For subcollimators 2-9, the model fits the observations quite well. For the finest subcollimator, the angular pitch is much smaller than the source size, and the amplitudes are essentially indeterminate. The residuals, shown by squares, are relatively small for all subcollimators above 1.<br />
<br />
Amplitude profiles such as this (and of course, phase profiles, which, for reasons of nugget space we do not show) are invaluable for diagnosing source structure of many kinds. <br />
<br />
== What use are RHESSI Visibilities? ==<br />
<br />
Aside from the advantages that RHESSI visibilities are a highly compact, device-independent representation of the data, there are several other advantages to creating them. For one, making maps can be greatly sped up by using highly optimized radio astronomy programs (see this [http://hesperia.gsfc.nasa.gov/~schmahl/NJIT/MEM_NJIT.pdf 18-page .pdf paper] for details). For another, one may reliably determine source sizes using a [http://hesperia.gsfc.nasa.gov/~schmahl/vis_fwdfit/fwdfit_tutorial.html visibility forward-fit routine]. Other (as yet) barely-exploited advantages are <br />
<br />
* Mapping in the 6.7 keV Fe line;<br />
* Mapping in the sum of nuclear lines;<br />
* Mapping in terms of photon energy, detected energy;<br />
* Mapping separately in nuclear lines and the continuum;<br />
* Improved pileup corrections for hard X-ray images<br />
* Enhancing statistical sensitivity by weighting<br />
* Improving iterative processing <br />
<br />
<br />
Another important current use is the self-calibration of phases. By means of such self-calibration it will become possible to utilize the 2nd and 3rd harmonics of the near-triangular RHESSI waveform (Figure 1). Up until now, we have used only the fundamental sinusoids, but the higher harmonics will provide more complete coverage of the Fourier plane for better imaging and higher resolution, as good as ~ 1 arcsec.<br />
<br />
== Ackowledgements ==<br />
<br />
The RHESSI software team, particularly Richard Schwartz and Rick Pernak, have helped bring us into a renaissance of visibilities. Without their continuing assistance, the team would still be in the dark ages of Fourier imaging.<br />
<br />
Biographical notes: Ed Schmahl and Gordon Hurford are senior RHESSI team members based at GSFC and UC Berkeley, respectively.<br />
<br />
[[Has article subject:: RHESSI]]<br />
[[Has article subject:: education]]<br />
[[Has observation by:: RHESSI| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_VisibilitiesRHESSI Visibilities2018-09-18T02:28:20Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = RHESSI Visibilities<br />
|title = Nugget Details<br />
|number = 39<br />
|first_author = Ed Schmahl<br />
|second_author = Gordon Hurford<br />
|publish_date = 11-Jun-2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::40]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::38]]}}<br />
}}<br />
<br />
== Visibilities à la Fourier ==<br />
<br />
RHESSI is a Fourier imager, which means that it obtains amplitudes and phases (a.k.a. "visibilities") of the X-ray image modulations. This particular choice for an imager [http://hesperia.gsfc.nasa.gov/~kim/Science.html#imaging_system design] was driven by many forces, not the least of which was the mantra "[http://www.astrosociety.org/pubs/mercury/9504/tenerelli.html faster, better, cheaper]" which put strong cost and size constraints on the [http://sunland.gsfc.nasa.gov/smex/ Small Explorer (SMEX)] satellites of the 1990s. Previous solar instruments that used Fourier techniques were [http://www.daviddarling.info/encyclopedia/H/Hinotori.html Hinotori], [http://www.astro.umd.edu/~ed/gifs/heidi_payload.gif HEIDI], and [http://gedas22.stelab.nagoya-u.ac.jp/HXT/ HXT], but RHESSI has unique new Fourier/visibility-based [http://hesperia.gsfc.nasa.gov/rhessidatacenter/imaging/imaging_concepts.pdf capabilities] that are only now beginning to be exploited.<br />
<br />
What does Fourier have to do with imaging? The 18th century mathematician [http://en.wikipedia.org/wiki/Joseph_Fourier Jean Baptiste Fourier] discovered the principle that any signal can be broken up into a [http://mathworld.wolfram.com/FourierSeries.html superposition of sines and cosines], or the equivalent amplitudes and phases, but he did not himself invent Fourier imaging. Applications of Fourier methods to imaging had to wait until the late 19th century when the amplitude and phase pairs of optical interferometry (see this [http://www.astro.lsa.umich.edu/~monnier/Publications/ROP2003_final.pdf 69-page .pdf document] for full detail) came into use, and these paired variables came to be known as visibilities, as in "fringe visibility". Later in the 20th century [http://www.geocities.com/capecanaveral/2309/page3.html radio] interferometry used the same entities, and now hard X-ray astronomers have joined the visibility game (see this [http://sprg.ssl.berkeley.edu/~ghurford/VisibilityGuide.pdf 8-page .pdf] document for details).<br />
<br />
== How do we create visibilities? ==<br />
<br />
To get actual sine and cosine components of hard X-ray images is difficult, but RHESSI uses a good approximation: triangular waveforms. In the ideal case - constant background, thin grids, and high photon rates - the modulated waveform generated by photons incident on one of RHESSI's sub-collimators has a triangular profile (Figure 1). The amplitude of the waveform is proportional to the intensity of the beam, and its phase and frequency depend on the direction of incidence. For complex sources, and over small rotation angles, the amplitude and phase of the waveform provide a direct measurement of a single Fourier component of the angular distribution of the source. Different Fourier components are measured at different rotation angles and with grids of different pitches.<br />
<br />
[[Image:39_triangle_waveform.png|600px|thumb|center|'''Figure 1''': Ideal hard X-ray grid response (solid curve), with its associated fundamental sinusoid (dashed curve).]]<br />
<br />
== From time tags to visibilities ==<br />
<br />
We illustrate the end-to-end process of computing visibilities for a particular flare interval, a single subcollimator, and a single energy band. The primary RHESSI data are photon time tags and pulse heights, which must be transformed in various ways before imaging is possible. RHESSI's time tags and pulse heights for a selected 100-ms interval are shown in Figure 2a. These time tags have not yet been binned in energy or time, so they are difficult to interpret by eye. Several gaps produced by cosmic-ray hits are present. <br />
<br />
[[Image:39_time-tagged_pulse_hts.png|600px|thumb|center|'''Figure 2a''': Time-tagged pulse heights]]<br />
<br />
A sequence of basic (linear) operations must be performed on these data to convert them to visibilities. The next step is to bin the data in energy and in uniform roll-angle bins, as shown in Figure 2b.<br />
<br />
[[Image:39_energy-binned_pulse_hts.png|600px|thumb|center|'''Figure 2b''': Energy-binned pulse heights]]<br />
<br />
The next step is to "histogram" the time tags into "roll bins," ie units of RHESSI rotational phase, as shown in Figure 2c for 3 rotations.<br />
<br />
[[Image:39_modulated counts.png|600px|thumb|center|'''Figure 2c''': RHESSI modulation profile. The green vertical lines mark the boundaries of spacecraft rotations, (units of 360 degrees) and the red tick marks indicate the roll bins selected for phase-bin stacking. Three sets of rollbins are identified by the colors blue, orange, and cyan.]]<br />
<br />
We now progress further towards constructing visibilities by binning the modulation cycles into roll bins, which are then binned again into 12 aspect phase bins. In this case there are 16 roll bins, three of which are outlined in color (blue, orange, and cyan) in each rotation. After phase binning, each of these rollbins can be co-added.<br />
<br />
[[Image:39_stacked_counts_3_rbins.png|600px|thumb|center|'''Figure 2d''': Stacked roll bins 3,4 & 5 taken from the three selected regions shown in corresponding colors in the previous figure.]]<br />
<br />
This process, called "[http://hesperia.gsfc.nasa.gov/~schmahl/the_stacker/xdoc_hsi_phz_stacker.png stacking]" increases the signal-to-noise ratio and provides a platform for computing a mean amplitude and phase. This generally improves the aspect phase coverage because the angular drift of the telescope causes the aspect phase to differ from one rotation to the next. After stacking we fit sinusoids to the profiles (red curves in Figure 2d).<br />
<br />
In order to calibrate the parameters generated by the fits, a number of corrections must be made. The grids have a finite thickness, which produces internal shadowing and modifies the triangular-like response; this also causes a decline of transmission as a function of the off-axis distance; the grid slats are not all perfectly aligned parallel to the optical axis, which leads to a "venetian blinding" effect, which makes even and odd half-rotations unequal; and there are detector-to-detector sensitivity differences. Fortunately, RHESSI was calibrated in many ways before launch, and the instrument has many self-calibration capabilities, so all of these effects can be calibrated out. Using the RHESSI calibrations, one may remove the instrumental dependence and convert the sinusoids to photon rates.<br />
<br />
The final steps in constructing visibilities are to output the amplitudes and phases of the fitted sinusoids, and compute error estimates based on photon statistics and instrumental systematics. Figure 2e shows the computed amplitudes in the upper panel, and the visibility phases in the lower panel. <br />
<br />
[[Image:39_amplitudes_and_phases.png|600px|thumb|center|'''Figure 2e''': Visibility amplitudes and phases derived from the stacked counts in Figure 2d. Roll bins 3,4 & 5 taken from the three selected regions are shown in the same colors as the previous figure.]]<br />
<br />
Note that the visibility amplitudes are in units of photon flux, unlike the stacked modulation profiles, which are in counts per bin. Since the calibrations have been applied, the visibilities are very nearly independent of the instrument. The importance of this instrumental independence is hard to overestimate. It makes it possible, for example, to use radio imaging techniques such as [http://hesperia.gsfc.nasa.gov/~schmahl/MEM_NJIT/mem_njit_tutorial.html MEM] for RHESSI imaging, and to do forward fitting ([http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=35 see earlier nugget]). We list several other applications of visibilities later.<br />
<br />
== RHESSI visibility amplitudes ==<br />
<br />
Above, we showed data for only a single subcollimator. But RHESSI's great imaging power results from its coverage of a wide range of spatial scales with its nine subcollimators. Here (Figure 3) we show the amplitude dependence as a function of both subcollimator (SC) and roll angle (PA). The X-axis is a linear combination of both variables in the form SC+PA/180. Thus, for example, the roll angles 0 to 180 degrees are shown for subcollimator 6 for X between 6.0 and 7.0. <br />
<br />
[[Image:39_ed.gif|600px|center|thumb|'''Figure 3''': Observed visibility amplitudes (blue crosses) for a flare interval as a function of subcollimator (SC=1-9) and position angle (PA=0-180°) of the grids. Each of the 9 vertical panels shows the amplitude as a function of PA for one subcollimator (labeled by red digits below the X axis). The red curve represents a model using two Gaussian sources, and the green squares show the residuals relative to the model. For a given subcollimator (6 and 7 are good examples), the amplitude rises and falls sinusoidally while the grids rotate from PA=0 to PA=180°. In general, such sinusoidal variation indicates an extended or double source.]]<br />
<br />
The observed amplitudes are indicated by blue crosses and their associated estimated standard errors are shown by vertical error bars. Note the variable amplitude as a function of roll angle for each subcollimator. This variation is caused by the "beating" of two flare sources against each other. Note also the gradual falloff of amplitude as the subcollimators become finer (towards smaller X). This falloff is the result of the finite sizes of the sources. When the angular pitch is smaller than the source size, the modulation amplitude is reduced (and the error bars become larger).<br />
<br />
The solid red curve shows the visibility computed for a model given by two Gaussian sources. For subcollimators 2-9, the model fits the observations quite well. For the finest subcollimator, the angular pitch is much smaller than the source size, and the amplitudes are essentially indeterminate. The residuals, shown by squares, are relatively small for all subcollimators above 1.<br />
<br />
Amplitude profiles such as this (and of course, phase profiles, which, for reasons of nugget space we do not show) are invaluable for diagnosing source structure of many kinds. <br />
<br />
== What use are RHESSI Visibilities? ==<br />
<br />
Aside from the advantages that RHESSI visibilities are a highly compact, device-independent representation of the data, there are several other advantages to creating them. For one, making maps can be greatly sped up by using highly optimized radio astronomy programs (see this [http://hesperia.gsfc.nasa.gov/~schmahl/NJIT/MEM_NJIT.pdf 18-page .pdf paper] for details). For another, one may reliably determine source sizes using a [http://hesperia.gsfc.nasa.gov/~schmahl/vis_fwdfit/fwdfit_tutorial.html visibility forward-fit routine]. Other (as yet) barely-exploited advantages are <br />
<br />
* Mapping in the 6.7 keV Fe line;<br />
* Mapping in the sum of nuclear lines;<br />
* Mapping in terms of photon energy, detected energy;<br />
* Mapping separately in nuclear lines and the continuum;<br />
* Improved pileup corrections for hard X-ray images<br />
* Enhancing statistical sensitivity by weighting<br />
* Improving iterative processing <br />
<br />
<br />
Another important current use is the self-calibration of phases. By means of such self-calibration it will become possible to utilize the 2nd and 3rd harmonics of the near-triangular RHESSI waveform (Figure 1). Up until now, we have used only the fundamental sinusoids, but the higher harmonics will provide more complete coverage of the Fourier plane for better imaging and higher resolution, as good as ~ 1 arcsec.<br />
<br />
== Ackowledgements ==<br />
<br />
The RHESSI software team, particularly Richard Schwartz and Rick Pernak, have helped bring us into a renaissance of visibilities. Without their continuing assistance, the team would still be in the dark ages of Fourier imaging.<br />
<br />
Biographical notes: Ed Schmahl and Gordon Hurford are senior RHESSI team members based at GSFC and UC Berkeley, respectively.<br />
<br />
[[Has article subject:: RHESSI]]<br />
[[Has article subject:: education]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Annealing_RHESSI_for_the_first_timeAnnealing RHESSI for the first time2018-09-18T02:28:01Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Annealing RHESSI for the first time<br />
|number=69<br />
|first_author=David Smith<br />
|publish_date=12 January 2008<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::70]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::68]]}}<br />
}}<br />
<br />
==Introduction==<br />
<br />
The RHESSI detectors are high-purity germanium semiconductor diodes, operated in a unique segmented manner (please see our <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/X-ray_and_gamma-ray_spectroscopy earlier Nugget] for more detailed information). All such detectors are sensitive to radiation damage by high-energy particles in space. In RHESSI's low-Earth orbit, the particles are those of the Van Allen belts, and the dosage can be readily predicted. Thus when RHESSI was launched we were aware that the detectors would have a finite useful lifetime. Gradually, as damage builds up, the detector resolution and effective volume drop.<br />
<br />
Annealing - accomplished by heating the detector to a relatively high temperature (say 100 C) and allowing it to soak at this elevated temperature (say for a week) can restore functionality. The RHESSI team had always planned this but became more confident as the other Ge detector array, that on the gamma-ray observatory [http://sci.esa.int/integral/ Integral], benefited from routine annealing operations.<br />
<br />
[[File:69f1.png|400px|thumb|center|<br />
Figure 1: Cross-section of one of RHESSI's cylindrical Ge detectors, showing the general structure of the electric fields in the detector volume. <br />
]]<br />
<br />
Figure 1 (above) shows the geometry of the RHESSI detectors, along with the pattern of the electric field within the Ge volume that sweeps out the charge left by a photon interaction. <br />
The Sun would be at the top, and the front segment is that part of the volume above the dashed line, the rear segment below. Separate anodes (contacts made in the hollow central region) separately collect the charge from photon interactions the two segments. <br />
Radiation damage to the crystal traps moving charges that are part of the signal from a gamma-ray detection, resulting in a partial loss of signal and poorer energy resolution. It also creates permanent charges throughout the detector volume, which can distort the imposed electric field and, when the damage is particularly severe, create dead volumes within the crystal.<br />
<br />
<br />
==The 2008 RHESSI anneal==<br />
<br />
After much preparation, the RHESSI team annealed the detectors in November for the first time, almost six years after launch. By this time the radiation damage had built up to such a degree that normal operation, especially for gamma-rays, was rapidly becoming difficult. The operation was worrisome, since the segmentation (see Figure 1) might have been destroyed by the anneal. The Integral detectors are not segmented and do not have this extra risk factor. That was the reason to wait so long for the first anneal, and to keep the temperature up only for a limited period of time (one week at 90 C). The result is shown in the figure below - a success roughly as expected, but not without surprises. The anneal did not restore the full resolution of the detectors, but that is of minimal importance. The main thing is that the sensitivity (volume and efficiency) have returned.<br />
<br />
[[File:69f2.png|600px|thumb|center|<br />
Figure 2: Response of the RHESSI rear segments at the 511-keV line of positron annihilation (a background feature). Green was the situation before the anneal, and red the situation afterwards. The black (initial) and blue (mid-2006) behavior shows that the anneal fully restored the detectors' effective volume if not the original energy resolution.<br />
The annealing also restored the front segments to good performance, as shown in Figure 3. Seven individual detectors are shown as separate curves for three solar flares observed in 2005 (left), then just before the anneal (middle), and at present after the anneal (right). Note the disorder and disagreement among the detectors just before (the middle panel) and the good agreement afterwards.<br />
]]<br />
<br />
[[File:69f3.png|600px|thumb|center|<br />
Figure 3: Spectra of three solar flares observed during normal operation (left), just before the anneal (center), and after the anneal (right). These spectra cover the soft X-ray range with a peak at about 7 keV. The radiation damage destroyed the agreement among the 9 detectors (central panel) making analysis very difficult, but the anneal has restored the agreement satisfactorily.<br />
]]<br />
<br />
==The future==<br />
<br />
With this successful annealing operation, RHESSI is ready for the energetic flares expected in Cycle 24, which is just beginning (see our earlier Nugget, in particular its Figure 3). We can do it again if need be, and in fact expect to do so as radiation damage again gradually builds up.<br />
<br />
==Biographical note==<br />
David Smith is a RHESSI team member at UC Santa Cruz.<br />
<br />
[[Has article subject:: RHESSI]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Optical_ImagesRHESSI Optical Images2018-09-18T02:27:41Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Nugget Details<br />
|number = 85<br />
|first_author = H. Jabran Zahid<br />
|second_author = Hugh Hudson<br />
|publish_date = 29 October 2008<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::86]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::84]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
RHESSI's sole scientific instrument is its array of X-ray and gamma-ray<br />
imagers, as described from many points of view in these Nuggets.<br />
To make this high-resolution imaging possible, though, even higher-resolution<br />
optical images are required.<br />
Only in this way can we follow the natural <br />
[http://en.wikipedia.org/wiki/Poinsot's_construction Poinsot motions] and thereby determine, photon by photon, the <br />
[http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=8 X-ray image] structure. <br />
<br />
This requirement dictated the invention of a solar aspect sensor (SAS),<br />
which consists of three simple lenses forming solar images on three<br />
linear CCDs, with 2048 x 1.73'' pixels.<br />
The set of three line images (plus independently obtained stellar aspect data) <br />
uniquely determines the instantaneous angular position of the RHESSI imager,<br />
to a resolution considerably finer than that of the X-ray images themselves.<br />
In this Nugget we show that these 1D images can be assembled into rather<br />
clean 2D images of the Sun in the 670-nm visible continuum.<br />
This unusual way of making images is of some scientific value, as will be<br />
related in future Nuggets. <br />
The merits of doing this are (a) that RHESSI is in space (no astronomical<br />
seeing whatsoever), and (b) RHESSI is rapidly rotating (radically <br />
different and more tractable <br />
[http://www.starlink.rl.ac.uk/star/docs/sun139.htx/node134.html flat-field] problem).<br />
<br />
== Some Data ==<br />
<br />
RHESSI normally provides about 1,000 line images per day, and they are<br />
spread around over all position angles (image azimuths), in the manner of<br />
sticks piling up on a map (somewhat resembling<br />
[http://mathworld.wolfram.com/BuffonsNeedleProblem.html Buffon's Needle Problem]). <br />
For each line image we determine the limb intercepts, thus defining a chord<br />
length, and map the pixel contents onto the 2D solar image according to the<br />
roll angle of the spacecraft motion.<br />
The spacecraft nutates as well as rotates, and so the actual sampling pattern<br />
is fairly complicated.<br />
The sampling is roughly sufficient to generate one clean image per day, <br />
but the longer the integration, the smoother the result.<br />
<br />
Figure 1 shows this in the form of a one-month average and a high-pass<br />
version of the same thing.<br />
Here "high pass" means high<br />
[http://en.wikipedia.org/wiki/Spatial_frequency spatial frequencies]; the <br />
[http://en.wikipedia.org/wiki/Sobel_operator Sobel filter]<br />
we have used strongly reduces the smoothly varying parts of the disk<br />
and thus enhances its fine structures - spots, faculae, and the limb itself.<br />
<br />
[[Image:RHESSI may image.gif|450px|thumb|center|'''Figure 1''': Left: image of the Sun for the month of May, 2003. Because of the sparse sampling, no features are immediately obvious. Right: high-pass filtering of the same image, emphasizing the location of the limb and also showing faint traces of active regions spread out by solar rotation.]]<br />
<br />
Note the noise patterns in the image on the right (click to enlarge). Other than the active-region traces themselves, there is also a non-uniform granularity of the images, plus excess noise in a small region near disk center. These properties result from the manner of RHESSI's rotation and nutation, but it is a somewhat subtle matter that we will not explain here.<br />
<br />
== Conclusions ==<br />
<br />
Of what use are images of this sort?<br />
We are just working on applications that will be described in future RHESSI <br />
Science Nuggets and, we hope, research papers.<br />
These will include the characterization of the facular limb-darkening<br />
function and the global temperature structure of the photosphere, both<br />
problems of outstanding importance.<br />
There is at least one better optical telescope viewing the Sun <br />
from a satellite, the 50-cm <br />
[http://www.youtube.com/watch?v=kxorDewEwbM SOT] on the<br />
[http://solarb.msfc.nasa.gov/ Hinode] spacecraft.<br />
The advantages of our RHESSI images include its whole-Sun imaging and <br />
superior photometric performance, thanks to the telescope rotation.<br />
We have not yet quantified the photometric precision sufficiently<br />
but believe that it will be extremely good, because each pixel<br />
redundantly observes many positions on the solar disk.<br />
In the meanwhile we are just pleased that this imaging techniques works<br />
so well and are presenting it because it is novel in this application.<br />
<br />
'''Biographical note''': Jabran Zahid is now a graduate student at the University of Hawai'i, but was a researcher at Berkeley when much of this work was done. Hugh Hudson is a senior RHESSI team member in Berkeley.<br />
<br />
[[Has observation by:: RHESSI SAS| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Chree_Analysis_for_FlaresChree Analysis for Flares2018-09-18T02:26:42Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Nugget Details<br />
|number = 98<br />
|first_author = Hugh Hudson<br />
|second_author = <br />
|publish_date = 30 March 2009<br />
|next_nugget = [[Cycle 24 - don't panic yet!]]<br />
|previous_nugget = [[High Temperatures in Active Regions]]<br />
}}<br />
<br />
<br />
<br />
== Introduction ==<br />
<br />
"Chree analysis," named after [http://www.nndb.com/people/824/000168320/ Charles Chree], is now normally called "superposed epoch analysis."<br />
It is a powerful but tricky device, used effectively by Chree himself but often the source of erroneous conclusions.<br />
In this kind of analysis one takes a data series that contains interesting features, and then co-adds short segments of the series using some other source of information to provide a registration index.<br />
Thus if one could barely hear [http://en.wikipedia.org/wiki/Big_Ben Big Ben] chime from some distant location, one could just make a recording and then add up the parts centered on the hour as deduced from any other accurate clock.<br />
Eventually the noise would average down, and Big Ben's chimes would become strong enough to hear.<br />
That is the simple view.<br />
Although Chree himself used his method to make many important discoveries about geomagnetism (see Ref. 1), <br />
an unwary scientist can easily blunder using the technique.<br />
<br />
== Flare occurrence ==<br />
<br />
Solar flares occur at odd intervals, but not randomly. <br />
There is an extensive literature discussing various aspects of this non-randomness, and Figure 1 (from a paper by Crosby et al.) illustrates one of them, namely the rigorous power-law pattern of flare magnitude.<br />
More energetic flares are less numerous, following the precise statistical relationship shown in the Figure.<br />
Below the weakest of these events one finds the [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=52 microflares] or even<br />
perhaps some of Parker's famous [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/The_Jakimiec_Track nanoflares].<br />
<br />
[[Image:98crosby.jpg|400px|thumb|center|'''Figure 1'': The rather nice power-law distribution found by Norma Crosby and collaborators <br />
to describe solar flare hard X-rays observed by the [http://heasarc.gsfc.nasa.gov/docs/heasarc/missions/solarmax.html HXRBS] instrument.]]<br />
<br />
The distribution function follows the power law <math>N(E) = AE^{-\gamma}</math>, where <math>A</math> (the normalization of the number of events per unit time) and <math>\gamma</math> (the power-law index) are constants.<br />
One interesting feature of this description (other than its amazing precision) is the flatness of the observed power law, i.e. the smallness of <math>\gamma</math> (it is 1.732 <math>\pm</math> 0.008; see Figure 1).<br />
As any math student can show, the integral of this distribution diverges, i.e. it incorrectly "predicts" an infinite flare power since the integral of E.N(E) diverges if <math>\gamma < 2</math>.<br />
This property is somewhat perplexing, since obviously the Sun does not shine by flares alone.<br />
But it is not very perplexing because there should be some physical reason to truncate this distribution, and indeed this truncation may have been observed<br />
already in an arcane piece of data: <sup>12</sup>C in tree rings ([http://en.wikipedia.org/wiki/Dendrochronology dendrochronology]), and fossil cosmic-ray tracks in the [http://en.wikipedia.org/wiki/Regolith lunar regolith] (Ref. 2)!<br />
<br />
== Signal or noise? ==<br />
<br />
Given a well-defined occurrence distribution, how much does one improve the signal-to-noise ratio (SNR) by including a range of flare magnitudes? <br />
What we would like to do is to improve on the very important observation of flare effects on the total solar irradiance (Ref. 3), as shown for the single best example in Figure 2 below.<br />
The important thing to notice is the fluctuation (solar in origin) against which the flare effect must be detected.<br />
Can we add many time series of this type, guided by knowing the epochs of flare occurrence, to defeat this noise?<br />
The epoch in this case (the maximum, as marked by the arrow in Figure 2) coincides closely with the hard X-ray peak time as observable by RHESSI, although unfortunately not in this case.<br />
<br />
[[Image:98flare_tsi.jpg|400px|thumb|center|'''Figure 2''': Observation of a flare signal in the [http://spot.colorado.edu/~koppg/TSI/ total solar irradiance] (ie, an apparent temporary increase of solar luminosity, as recorded by the [http://lasp.colorado.edu/sorce/index.htm SORCE] satellite. <br />
Note the heavily suppressed zero of the flux scale (W/m<sup>2</sup>).<br />
Also note the fluctuating signals before the flare; this is the background fluctuation (signal for some, but noise for this purpose);<br />
this noise greatly reduces the precision of the flare observation.]]<br />
<br />
Clearly one would benefit from a Chree analysis of (say) a set of flares all of the same magnitude.<br />
Roughly the SNR would increase by about N<sup>1/2</sup> as N increases.<br />
But if one simply waits for flares to occur and takes them all, regardless of magnitude, would that help?<br />
The answer turns out to depend on the value of the power-law index.<br />
For a steep power law <math>\gamma \gg 2 </math> the actual value does not matter, and one gains the full N<sup>1/2</sup> factor by including flares of all magnitudes.<br />
But for a slope less than the critical value 2, it does matter. <br />
Then it can be shown that the SNR grows much more slowly with magnitude range. <br />
For R = M<sub>max</sub>/M<sub>min</sub> (R the ratio of the strongest and weakest events in the sample), we find that the SNR increases only as R<sup>2-&gamma;</sup>, much more slowly than the increase in number of events.<br />
The closer the power-law index comes to 2, the slower the gain via a Chree-type analysis.<br />
<br />
== Conclusions ==<br />
<br />
Even a simple-seeming and intuitively correct analysis technique can have tricky consequences.<br />
The Chree analysis does work in this case (see Ref. 3), but not as well as one would like. <br />
For a given SNR for the detection of a single X10-class flare, the extension to the M1 level (a factor of 100 in magnitude) would correspond to an increase in SNR of 3.6 (using the slope from Figure 1).<br />
While this is certainly not to be sneezed at - it might be just the factor needed to learn something really important about flare energetics - the weak dependence suggests that it is a difficult analysis.<br />
<br />
== References ==<br />
<br />
1. Chapman, S. (1941), ''Charles Chree and his work on geomagnetism'' [http://adsabs.harvard.edu/abs/1941PPS....53..629C 1], [http://adsabs.harvard.edu/abs/1941PPS....53..635C 2], [http://adsabs.harvard.edu/abs/1941PPS....53..650C 3] (see especially part 2)<br />
<br />
2. Reedy, R.C., Arnold, J.R., and Lal, D. (1983), ''Cosmic-ray record in solar system matter'' [http://adsabs.harvard.edu/abs/1983Sci...219..127R]<br />
<br />
3. Kretzschmar, M. et al. (2010), ''The effect of flares on total solar irradiance'' [http://adsabs.harvard.edu/abs/2010NatPh...6..690K]<br />
<br />
[[Has article subject:: education]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_-_Concept_to_FruitionRHESSI - Concept to Fruition2018-09-18T02:26:10Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI - Concept to Fruition<br />
|number = 100<br />
|first_author = Brian Dennis<br />
|second_author = Bob Lin<br />
|publish_date = 30 April 2009<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::101]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::99]]}}<br />
}}<br />
<br />
== Introduction ==<br />
RHESSI (the [http://en.wikipedia.org/wiki/Reuven_Ramaty Reuven Ramaty] High-Energy Solar Spectroscopic Imager) had a long and tortured gestation period from the initial concept perhaps as long ago as the 1970's to its final selection as a Small Explorer (SMEX) mission in 1997 and its subsequent delayed launch in 2002. We take the opportunity of [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Science_Nuggets RHESSI Science Nugget] number 100 to summarize the many steps along the way.<br />
The original observational goals that ultimately led to RHESSI were imaging and spectroscopy of solar flares in hard X-rays and gamma rays, with high angular resolution and fine energy resolution. <br />
Earlier instruments had tackled one or the other of these objectives but it was only with RHESSI that both were achieved to provide imaging spectroscopy over a broad energy range. <br />
This Nugget describes how many lines of study, development, and proposal eventually led to RHESSI.<br />
<br />
== Hard X-ray Imaging ==<br />
<br />
The [http://heasarc.gsfc.nasa.gov/docs/heasarc/missions/solarmax.html Solar Maximum Mission], launched in 1980, carried the Hard X-ray Imaging Spectrometer (HXIS), the first instrument capable of imaging in X-rays up to 30 keV with 8 arcsecond pixels. <br />
''Hinotori'' followed soon thereafter using rotating modulation collimators to achieve similar angular resolution but greater sensitive area. <br />
And then of course, the Hard X-Ray Telescope (HXT) on [http://en.wikipedia.org/wiki/Yohkoh Yohkoh] in the 1990's improved on the angular resolution and extended the energy range to close to 100 keV.<br />
<br />
In the meanwhile, many other ideas appeared, but could not be realized for various reasons.<br />
The ambitious Pinhole Occulter Facility (P/OF - pronounced "poff") was proposed in the 1980s (see Figure 1); it would have used a boom that could be extended to 50 m to achieve sub-arcsecond resolution into the gamma-ray range. <br />
It was initially proposed for the Space Shuttle (see Figure 1) and later for the Space Station but was never selected for funding, even though a [http://www2.jpl.nasa.gov/srtm/mast.html flexible boom] was at one point deployed from the Shuttle.<br />
<br />
[[Image:POF.jpg|200px|thumb|left|'''Figure 1''': The Pinhole/Occulter Facility (P/OF) as conceived in 1985 for the Shuttle.]]<br />
<br />
One instrument that did make it (almost) to space was the [http://astrophysics.gsfc.nasa.gov/balloon/ balloon-borne] High Energy Imaging DevIce (HEIDI). It tested out the rotating modulation collimator and the solar aspect system concepts that are used on RHESSI. <br />
HEIDI was built at [http://www.nasa.gov/centers/goddard/home/index.html Goddard] and flown once from [http://www.csbf.nasa.gov/ Texas] in 1993.<br />
<br />
== Hard X-ray and Gamma-Ray Spectroscopy ==<br />
<br />
The first gamma-ray spectroscopy of flares was achieved on [http://heasarc.gsfc.nasa.gov/docs/heasarc/missions/oso3.html OSO-3] using sodium iodide NaI(Tl) scintillation detectors, and similar observations were made from [http://heasarc.gsfc.nasa.gov/docs/heao1/heao1.html HEA0-1] and [http://heasarc.gsfc.nasa.gov/docs/heasarc/missions/solarmax.html SMM] in the 1980s. <br />
All of these had low spectral resolution, insufficient to resolve the physically important spectral scales.<br />
The instrument that evolved into RHESSI's high spectral resolution capability was the balloon-borne HIgh REsolution X-ray spectrometer (HIREX) built at Berkeley with cooled planar germanium detectors. <br />
This was flown several times in the U.S. and from the Antarctic and made seminal observations of the "super-hot component," solar microflares, and galactic center X-ray sources. A later hard X-ray and gamma-ray spectrometer using large, segmented, coaxial, germanium detectors (see the [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 relevant nugget]) was built in the 1990s as the HIgh REsolution Gamma-ray and hard X-ray Spectrometer (HIREGS).<br />
It has flown several times on long-duration balloons from the Antarctic.<br />
<br />
== Combined Imaging and Spectroscopy ==<br />
<br />
An early audacious idea was to put P/OF on SMM during a second repair mission in the mid-1980s but this was quickly rejected as too costly. However, it did lead rather directly to the Max ’91 program and the Solar High-Energy Astrophysical Plasmas Explorer (SHAPE; see Figures 2 and 3) proposal that was submitted in 1986. SHAPE still had a separate imager and spectrometer on a 3-axis stabilized spacecraft (i.e, not spin-stabilized like RHESSI). The Gamma-Ray Imaging Device (GRID) would have had 34 bi-grid modulation collimators with the grids separated by 6.7 m to give arcsecond imaging; the High-resolution Gamma-Ray and Neutron Spectrometer (HIGRANS) would have had 12 dual-segment high-purity germanium detectors to give keV energy resolution, and a 5-cm thick BGO shield for low background. <br />
<br />
[[Image:GRID.jpg|200px|thumb|right|'''Figure 2''': The GRID instrument from the SHAPE proposal.]]<br />
<br />
[[Image:HIGRANS.jpg|200px|thumb|right|'''Figure 3''': The HIGRANS instrument from the SHAPE proposal.]]<br />
<br />
Unfortunately, the [http://en.wikipedia.org/wiki/Space_Shuttle_Challenger_disaster Challenger disaster] occurred just before the SHAPE proposal was submitted, and all new missions were put on hold. By the time the review panel made the selection it was deemed to be too late for the solar maximum of 1990 and the proposal was rejected. However, HEIDI and HIREGS grew out of this effort and provided the basis for the High Energy Solar Physics mission (HESP - a name that was stolen from the Japanese, who used it as the original name for ''Yohkoh''). HESP, shown in Figure 4, was proposed in 1991 by a science study group of the same name. <br />
HESP was a proposed medium-class satellite, unfortunately never flown.<br />
However it was the first time that a combined imager/spectrometer was proposed as a single instrument with the same basic elements as in the final HESSI proposal, a major step towards the first [http://en.wikipedia.org/wiki/Imaging_spectroscopy imaging spectroscopy] in the hard X-ray/gamma-ray spectral domain.<br />
<br />
[[Image:HESP.jpg|200px|thumb|left|'''Figure 4''': The High Energy Solar Physics mission (HESP) proposed in 1991 by a science study group with the same name.]]<br />
<br />
HESP was followed by the High Energy Solar Imager (HESI) in 1995, the MIDEX version of HESSI proposed in 1995 (with the second "S" added to acknowledge the importance of spectroscopy), and finally the successful SMEX proposal in 1997.<br />
<br />
== Selection to Launch ==<br />
<br />
After selection, HESSI quickly became reality with a compressed fabrication and qualification schedule made possible by the PI-controlled mode of operation. <br />
Most of the work was done at the [http://www.ssl.berkeley.edu/ Space Sciences Lab] in Berkeley, and at NASA's [http://hesperia.gsfc.nasa.gov/hessi/ Goddard Space Flight Center]) with major Swiss contributions from the [http://www.psi.ch/index_e.shtml Paul Scherrer Institut] and [http://www.ethz.ch/ ETH].<br />
<br />
After being ahead of schedule and within budget for an auspiciously planned 4 July 2000 launch, disaster struck during the vibration tests at JPL in December 1999. <br />
A malfunction of the "shake table" resulted in subjecting the entire satellite to [http://en.wikipedia.org/wiki/G-force G-forces] that were much higher than the design limits. <br />
The solar panels and two of the three mounts holding the telescope to the spacecraft were broken, as well as some other parts of the instrument.<br />
We had to remount and align many of the grids, replace the cooler with a backup, and start the qualification process from the beginning. <br />
Despite this ''shattering'' setback, we were ready for launch just six months late in December 2000. <br />
After some delay, the launch date was set for 7 June 2001 but another disaster beyond our control struck just 6 days earlier - a [http://en.wikipedia.org/wiki/Pegasus_(rocket) Pegasus rocket] of the same type that was to launch HESSI had gone out of control in a test flight of NASA's [http://www.nasa.gov/missions/research/x43-main.html X-43A] ramjet-powered test vehicle on June 2, 2001. <br />
Fixing the Pegasus problem pushed the HESSI launch date into 2002. <br />
<br />
A final scare occurred over the Atlantic Ocean as the [http://www.dfrc.nasa.gov/Gallery/Movie/Pegasus/HTML/EM-0024-02.html plane] carrying the Pegasus rocket was circling over the "drop zone." <br />
An "open mike" in the plane's communication system caused the pilot to make a second [http://en.wikipedia.org/wiki/Go-around go-around], taking an extra 20 minutes before the final drop and leading to an unbelievable feeling of relief when the Pegasus rocket ignited. <br />
After a perfect flight, RHESSI was put into its nominal 600-km circular orbit on 5 February 2002, some 20 months after the originally planned launch date.<br />
<br />
== Operations ==<br />
<br />
Once in orbit, things went smoothly, proving that space is much more benign than anywhere on Earth. The mechanical [http://auto.howstuffworks.com/stirling-engine1.htm Stirling-cycle] cryocooler, thought to have been the highest risk item, performed flawlessly cooling the detectors down to their ~80 K operating temperature in about a week. <br />
We note here that after more than seven years, this refrigerator is still cooling the detectors to below 100K.<br />
The spacecraft was set spinning at 15 RPM and the first flare was successfully detected (images and spectra!) on 12 Feb. 2002. Despite the delayed launch, there was plenty of solar activity left in the cycle and over 40,000 flares are now included in the catalog of recorded events. <br />
With these observations, RHESSI has provided the best view to date of the hard X-ray bremsstrahlung emissions, while at the same time doing remarkable new things such as gamma-ray imaging and high-resolution spectroscopy. These Nuggets (both the [http://sprg.ssl.berkeley.edu/~tohban/nuggets/ old series] and our current [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Science_Nuggets Wiki series]) provide many glimpses into these discoveries.<br />
<br />
== RHESSI's proposal history ==<br />
<br />
* 1980: First HIREX solar balloon flight<br />
* 1980: <font color = 'red'> Launch of the Solar Maximum Mission </font><br />
* 1981: <font color = 'red'> Launch of ''Hinotori'' </font><br />
* 1983: P/OF study (P/OF = "Pinhole Occulter Facility"; see Figure 1)<br />
* 1986: P/OF study<br />
* 1986: Max '91 study <br />
* 1986: SHAPE proposal (see Figures 2 & 3)<br />
* 1991: <font color = 'red'> Launch of ''Yohkoh'' </font><br />
* ~1991 - 1993: HiREGS balloon flights<br />
* 1993: HEIDI balloon flight (led to RHESSI's aspect sensor [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=19 (1)], [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Optical_Images (2)])<br />
<br />
* 1995: HESP study (see Figure 4)<br />
* 1995: HESI study<br />
* 1995: MIDEX proposal<br />
* 1997: SMEX proposal (this became RHESSI!)<br />
* 2002: <font color = 'red'> Launch of RHESSI </font><br />
<br />
== Conclusions ==<br />
<br />
The saddest part of the whole adventure was that Reuven Ramaty, who some have called one of the fathers of solar gamma-ray astronomy, died just months before RHESSI was launched. He had been a key figure in all of the planning and proposal efforts and it was fitting that HESSI be renamed after him. It is now known as RHESSI in his honor.<br />
<br />
Now, after over seven years in orbit, RHESSI is beginning to show its age a little bit. The cryocooler is slowly losing efficiency so that the germanium detectors are now operating at somewhat higher temperatures than optimum. The detectors also naturally degrade in the intense radiation bombardment of space, reducing their sensitive volume, sensitivity, and spectral resolution. These effects can be mitigated by annealing the detectors at close to 100 degrees Centigrade for about a week. This process was carried out for the first time in November 2007 but needs to be repeated [http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Cycle_24_-_don%27t_panic_yet%21 when solar activity resumes]. We expect that RHESSI will then continue to make important X-ray and gamma-ray imaging spectroscopy observations of solar flares well into the future. <br />
<br />
New results from old RHESSI data are frequently appearing even now during activity minimum.<br />
Future RHESSI observations will benefit from contemporaneous measurements made at other wavelengths with the many powerful advanced instruments that are already operational on Hinode, STEREO, CORONAS-Photon, etc., or will soon become so - e.g., the Solar Dynamics Observatory. Consequently, we expect plenty of exciting new material for these Nuggets and look forward to the next 100 in the series.<br />
<br />
[[Has article subject:: RHESSI]]<br />
[[Has article subject:: history]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/How_does_RHESSI_make_images%3FHow does RHESSI make images?2018-09-18T02:24:05Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=How does RHESSI make images?<br />
|number=8<br />
|first_author=Gordon Hurford<br />
|second_author= Steven Christe<br />
|publish_date=29 August 2005<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::7]]}}<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::9]]}}<br />
}}<br />
<br />
==How does RHESSI make images?==<br />
<br />
An important observational goal of [http://hesperia.gsfc.nasa.gov/hessi/| RHESSI] is to make sharp images of solar flares in X-rays and gamma-rays at energies from 3 keV to 15 MeV. Over much of this energy range, there is no known material that can effectively reflect or refract X-rays. How do we make sharp images if we cannot use lenses or mirrors?<br />
<br />
The answer is to selectively block the X-ray photons. If this is done in a way that depends on which direction the photons are coming from, then we can get the information needed to make an image.In practice we rapidly block and unblock incoming photons and use the pattern of rapid time-variations in the observed signal to tell us the exact direction to the X-ray source(s), their size, shape, strength, etc - in short, all the information needed to make an image.<br />
<br />
The device that does this is called a 'modulation collimator,' an invention attributed to [http://www.physicstoday.org/pt/vol-54/iss-9/p74.html |M. Oda]. Such a collimator consists of a pair of widely spaced grids located in front of a good-sized X-ray detector (needed for sensitivity). The grids have large numbers of parallel slits and slats made from a heavy material (tungsten or molybdenum) that is effective at stopping X-rays. The detector, located behind the grids, tells us the exact arrival time and energy of those X-rays that succeeded in getting through both sets of slits. Of course, it doesn't know about the X-rays that got stopped along the way. This is illustrated in Figure 1.<br />
<br />
Figure 1: An illustration of how the grids (slats seen end-on in red) block the X-rays seen by the detector (blue rectangle) below the grids. The plot at the bottom plots the number of detected X-rays against time. Click the plot to see what happens as the incident angle is changed with time. Note: The animation may play slowly the first time through.<br />
<br />
In the figure, the slats in the top grid cast an X-ray shadow onto the rear grid where a fraction between 0 or 100% of the remaining X-rays will reach the detector. This fraction depends on whether the shadow falls on the slits or the slats in the rear grid. To go from one extreme to the other requires a only a very small change in angle and so provides the basis for making very sharp images. The change in angle (in radians) is just equal to the ratio of the slit-to-slat distance to the separation of the grids. If two grids like those in Figure 2 are mounted 5 feet apart (as they are on RHESSI) that angle is only 2.3 arcseconds, or about 1/850th of the diameter of the Sun!<br />
<br />
Figure 2: A photo of one of RHESSI's grids. This one is made by stacking about 60 sheets of molybdenum to make a single 1 mm thick grid with 9 cm across. There are 2646 slats here; they are too fine to see in the photo (where they go from bottom left to upper right). However they can be seen as the horizontal bars in the inset. Each slat is separated from the next by 1/30th of a millimeter, about the thickness of a human hair.<br />
The movie shows that if we were able to change the incident direction of the X-rays, the number of detected X-rays could be made to vary quite rapidly, giving a distinctive modulated signal that can be measured. But how do we change the direction from which the X-rays come?<br />
<br />
That's the easy part. The collimator (grids and detector) are mounted on a rotating spacecraft. (RHESSI is pointed towards the Sun and spins at about 15 rpm.) From the point of view of someone (or something) fixed on the spacecraft, the stars, sunspots and anything else in the sky (like X-ray stars or gamma-ray bursters, for example) appear to move in a circle about the center of the Sun as the spacecraft rotates. (It's the same idea someone standing on the rotating Earth seeing the stars slowly move in a circle about the north pole. For RHESSI, it takes 4 seconds to go around the circle instead of 24 hours for the Earth. Looking along the slats as in Figure 1, the incident direction of the X-rays appears to move back and forth and so we a get a rapidly modulated signal (called a modulation pattern).<br />
<br />
How do modulation patterns like those in Figure 1 tell us source characteristics such as strength, location, size, shape, etc that we need to make an image? That can be understood by seeing how the modulation pattern changes when the source is changed in various ways. Below we reproduce the original plot in Figure 1 along with the solar source which would produce the modulation. <br />
<br />
The following plots show what the modulation pattern looks like if we make one change at a time in the source. Click on the plot to compare it to the original case.<br />
If the source were weaker, then the pattern just gets lower in amplitude but otherwise doesn't change its shape. <br />
<br />
If the source were located further from the axis of rotation, then there must be more cycles during each 4s rotation. <br />
<br />
If the source were located the same distance from the axis of rotation but at a different angle compared to North, then the pattern is shifted in time, but otherwise stays the same. <br />
<br />
If the source were larger in diameter but still put out the same number of X-rays, then the average number of detected X-rays would stay the same, but the modulation pattern gets smeared out and the variations get smaller. <br />
<br />
Now that we've seen how each characteristic of the source (size, strength, location, etc) has a distinctive effect on the measured modulation pattern, we get to the tricky part. In practice we have to work backwords! <br />
<br />
Instead of taking a known source and deducing what modulation pattern it would give, we must start with an observed modulation pattern and figure out what the original source looked like. Using the examples as a guide, this can be fairly straightforward if there is only one source and it has a simple shape. Real solar flares aren't so cooperative, however - there are often several sources at a time, some of which might have more complicated shapes and all of which combine together to give just one observed modulation pattern. It's a nice mathematical puzzle to work backwards and figure out the image that would give observed modulation patterns that look like this. <br />
<br />
How the RHESSI software solves this problem to generate images routinely will be the topic of a future nugget. Stay tuned.<br />
<br />
'''Biographical note''': Gordon Hurford is a scientist at the Space Sciences Laboratory.<br />
<br />
[[Category:Nugget needs figures]]<br />
<br />
[[Has article subject:: RHESSI]]<br />
[[Has article subject:: education]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/HESSI_and_Type_III_Radio_BurstsHESSI and Type III Radio Bursts2018-09-18T02:22:15Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = RHESSI and Type III Radio Bursts<br />
|number = 2<br />
|first_author = Steven Christe<br />
|publish_date = 9 April 2005<br />
|next_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::3]]}}<br />
|previous_nugget = {{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::1]]}}<br />
}}<br />
<br />
==Introduction==<br />
The RHESSI hard X-ray spectra and images, which have much greater sensitivity than previously available, may show for the first time the direct X-radiation produced by currently ill-understood coronal particle beams. Type III radio bursts have been observed by radio astronomers with such beams since the 1950s, but never heretofore been identified directly.<br />
<br />
A standard RHESSI image of a solar flare usually shows two features: at high energies, one or more compact sources, and at lower energies, a more diffuse structure a shape suggesting a magnetic loop. We identify these, respectively, as the footpoints of flaring loops and by the coronal bodies of these loops. The loops themselves, after cooling down from X-ray temperatures of 10-20 x 10^6 K, are clearly seen by [http://umbra.nascom.nasa.gov/eit/| EIT], [http://vestige.lmsal.com/TRACE/| TRACE], or in ground-based observations such as in [http://gcmd.nasa.gov/records/GCMD_FE00534.html| those] in [http://scienceworld.wolfram.com/physics/BalmerLines.html| H-alpha]. RHESSI can also distinguish these different sources spectroscopically, for example by use of the spectrograms reported in an earlier RHESSI science nugget. The low-energy sources are caused by violently heated gas. The high-energy footpoint sources result from energy losses by non-thermal electrons. In the standard model, as shown for example in this [http://solarmuri.ssl.berkeley.edu/~hhudson/cartoons/thepages/Cliver.html| cartoon], these electrons originate in the solar corona. By a mechanism not currently understood, they gain high energies and become "mildly relativistic" as they follow the magnetic field of the flare loops and crash into the high densities of the [http://csep10.phys.utk.edu/astr162/lect/sun/chromosphere.html| chromosphere] and [http://csep10.phys.utk.edu/astr162/lect/sun/photosphere.html| photosphere]. In addition to emitting the hard (10-100 keV) X-rays of our footpoint sources, these electrons also carry enough energy to cause the heating of the hot material in the coronal loop sources (1-2 keV).<br />
<br />
Figure 1 shows RHESSI images across the hard X-ray spectrum that illustrate this source structure: <br />
<br />
[[File:Nugget2 Feb20 images.png|700px|thumb|center|Figure 1:RHESSI images (2002 feb 20, 11:06:00-11:06:40) of a typical event, with photon energy increasing from left to right as shown. The harder images (at the right) begin to show artifacts (noise) resulting from photon counting statistics. Brighter colors represent larger flux. <br />
Image courtesy S. Krucker.]]<br />
<br />
As one can see (click image to enlarge; each image frame is 64 arc sec on a side) a simple two-footpoint structure becomes very clear at higher hard X-ray energies. Because each end of a coronal loop has a footpoint in the lower atmosphere, this double-source pattern is characteristic. Barring some technical questions (such as the all-important one of what accelerated the particles in the first place!), this picture of a solar flare is very robust, yet it is incomplete since it is known that accelerated particles sometimes escape the solar corona instead of going down the loop footpoints.<br />
<br />
==Radio type III bursts==<br />
<br />
The presence of escaping particles is known because they emit radio waves in a characteristic pattern: the emission frequency decreases rapidly with time, from for example 1 GHz down to the lowest frequencies observable from the Earth's surface (about 20 MHz). Below this frequency the [http://en.wikipedia.org/wiki/Ionosphere| Earth's ionosphere] normally absorbs radio waves, so prior to the space age this was effectively the lowest frequency available for radio astronomers to observe. At still lower frequencies the type III bursts often continue nevertheless, and are sometimes detected at Earth by satellites such as Wind or ACE. Just as RHESSI makes use of spectrograms, so do radio observatories: energy is replaced by frequency. Since the radio waves they emit is related to the density (there is a strong tendency for emission at the [http://scienceworld.wolfram.com/physics/PlasmaFrequency.html| plasma frequency]), the escaping particles show a particular radio signature where the emission drifts from high frequency (or density) to low frequency (or density). This kind of radio emission is known as a type III radio burst. These same particles should also produce X-rays similar to those seen at the loop footpoints. <br />
<br />
Here is a suggestive set of events. Has RHESSI observed the X-rays from the particles which generated the type III emission or are there other accelerated particles which have travelled down to the footpoint? <br />
<br />
In this figure please note the event at 14:30 especially: the radio spectrogram (bottom) shows a type III burst, and the X-ray spectrogram (middle panel) shows a hard X-ray spectrogram.<br />
<br />
[[File:nugget2_Christe_f1.gif|700px|thumb|center|Figure 2:Overview: Panel 1,2-GOES Level, Panel 3 - RHESSI spectrogram, Panel 4 - RHESSI Hard and Soft Channel, Panel 5 - WIND/WAVES Spectrogram]]<br />
<br />
==Conclusions==<br />
<br />
Ideally we can use RHESSI to track non-thermal electrons as they travel around, anywhere in the solar corona; practically we are limited by many factors (such as instrumental and solar background emission). Thus at present we don't know to what extent the above example represents an actual X-ray detection of the type III burst's escaping electrons as they actually escape. But we hope to have clear evidence soon, as RHESSI software and calibrations improve.<br />
<br />
[[Has article subject:: Type III radio bursts]]<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has observation by:: GOES XRS| ]]<br />
[[Has observation by:: WIND WAVES| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Novel_X-ray_SpectrogramsNovel X-ray Spectrograms2018-09-18T02:19:56Z<p>Schriste: </p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Novel X-ray Spectrograms<br />
|number = 1<br />
|first_author = Hugh Hudson<br />
|second_author = Sam Krucker<br />
|publish_date = 2005 March 21<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::2]]}}<br />
|previous_nugget=NA<br />
}}<br />
<br />
== Introduction==<br />
<br />
This is the first of a weekly series of RHESSI science nuggets, so it is an opportunity for an introduction to a basic tool that we use. [http://hesperia.gsfc.nasa.gov/hessi/| RHESSI] is a small NASA-sponsored space observatory, which carries the first [http://hesperia.gsfc.nasa.gov/hessi/hessispec.htm| high-resolution germanium detectors] capable of measuring X-rays and γ-ray spectra from the whole sky. It can also make high-resolution images of solar X-ray and γ-ray sources, as we will see in many future nuggets. RHESSI, to complete the introduction, is named for [http://www.physicstoday.org/pt/vol-54/iss-11/p80.html| Reuven Ramaty], a NASA scientist and a pioneer theorist of the evolving field of γ-ray astronomy.<br />
<br />
This science nugget introduces one simple concept that rapidly became a vital tool in RHESSI data analysis: the spectrogram (a.k.a. dynamic spectrum). We explain this below and explain its history as a research tool in astrophysics.<br />
<br />
==Hard X-ray Spectrograms==<br />
<br />
A spectrogram, as we use the term, is an image displaying time on the X-axis and the spectral distribution on the Y-axis, as in the example here. It is an image in which one coordinate is time, and the other photon energy. The brightness (or color level) then represents the spectral flux, usually as a logarithm to make the dynamic range larger in the reproduction. This survey tool makes sense for the first time (for solar X-rays) with the RHESSI data, since they are the first with high spectral resolution, and one can view the entire database this way via the [http://sprg.ssl.berkeley.edu/~tohban/browser/|RHESSI data browser].<br />
<br />
[[File: Article1_spectrogram.jpg|center|thumb|400px]]<br />
<br />
This of course is not a typical example; it is RHESSI's first good observation of a major solar flare, that of April 21, 2002. The time axis covers about three hours and has black gaps during orbit night (about 30 minutes during each 96-minute orbit). The vertical axis shows the hard X-ray spectrum, 3-200 keV, and there are many features (some not relevant to the flare) worth noting. The color table has red at the bright end, and pale blue/purple at the faint end. Click on the image to get a larger version. We note the following:<br />
<br />
* The flare itself, reaching a peak at the end of the first orbit: a broad diffuse horizontal band, peaking usually at 10-15 keV.<br />
* A series of sharply-defined vertical artifacts extending to the highest energies as dark regions; these are shutter-out intervals when the detectors saturate;<br />
* A series of hard X-ray spikes extending above 100 keV; these are the solar hard X-rays of the impulsive phase; and<br />
* An initial shutter transition early in the event, when the spectral peak energy moves from 5 keV first to about 7 keV, then to about 10 keV as the shutter closes.<br />
* Much of the behavior in this histogram is dictated by the autonomous operation of the RHESSI system of attenuating shutters, a simple but powerful innovation.<br />
<br />
Much of the behavior in this histogram is dictated by the autonomous operation of the RHESSI system of [http://hesperia.gsfc.nasa.gov/hessi/spacecraft.htm|attenuating shutters], a simple but powerful innovation.<br />
<br />
== Another good example, with an application ==<br />
<br />
In the example on the left, we have added a couple of things (again, click to enlarge). At the top there are now time series showing the GOES X-ray photometers and selected RHESSI counting rates, mostly near background levels. Microflares appear in the RHESSI spectrogram, second panel from the bottom. At the very bottom is a radio spectrogram from the [http://lep694.gsfc.nasa.gov/waves/waves.html|WIND] spacecraft, showing the passage of solar energetic electrons out through the corona from the microflares. After the gap for orbit night (10:30-11:00 UT) one sees other examples with different X-ray/radio relationships. Finally, RHESSI background features appear at various places, notably the pesky Ge background emission line at about 10 keV, as horizontal lines.<br />
<br />
[[File:Article1_spectrogram_ii.jpg|thumb|center|400px]]<br />
<br />
== Conclusion ==<br />
<br />
The spectrogram view of spectra arrayed in a time series, used in other fields and branches of astrophysics, now has a good application for solar hard X-rays thanks to RHESSI. It reveals a wealth of structure and allows one to compare morphological features of this structure between one kind of observation and another (sample scientist's comment: "gee whiz"). In particular the second example above shows how very sensitive RHESSI is, with excellent fine structure visible in the spectrogram even for tiny microflares.<br />
<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has article subject:: RHESSI]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Relative_and_(maybe)_Absolute_RHESSI_Detector_Efficiency:_2002-2008Relative and (maybe) Absolute RHESSI Detector Efficiency: 2002-20082018-09-18T02:18:35Z<p>Schriste: /* References */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Nugget Details<br />
|number = 117<br />
|first_author = Jim McTiernan<br />
|publish_date = 22 December 2009<br />
|next_nugget = [[Cycle 24 has begun]]<br />
|previous_nugget = [[A tiny white-light flare]]<br />
}}<br />
<br />
== Introduction ==<br />
<br />
[http://hesperia.gsfc.nasa.gov/hessi/ RHESSI] has now been observing for nearly eight years. How have the [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=27 detector] sensitivities changed since it was launched? To characterize this important detector property we use a method and software devised by Brian Dennis and Kim Tolbert. This is based on fitting a simple thermal spectrum, for which there are two free parameters, to a given flare observation near the peak time. The parameters are temperature, T, and emission measure, [http://en.wikipedia.org/wiki/Spectroscopy EM]. We do this first for detector 1. Next, thermal components are fit to the spectra for the other eight detectors, with T fixed at the value found for detector 1, T<sub>1</sub>. This gives us a measure of the relative sensitivity of the detectors.<br />
<br />
== The Flare Sample ==<br />
<br />
We want to do this calculation for as many flares as possible. Here we chose every flare on the [http://http://hesperia.gsfc.nasa.gov/hessidata/dbase/hessi_flare_list.txt RHESSI Flare List] which was observed with no attenuators, no particle events, no data gaps, with the spacecraft at low [http://en.wikipedia.org/wiki/Geomagnetic_latitude geomagnetic latitude], and at least 5 minutes from the [http://en.wikipedia.org/wiki/South_Atlantic_Anomaly SAA]. <br />
These are all technical matters details that might bias the results of the cross-calibration and thus make it less precise.<br />
In addition to the flare interval being "clean" in this manner, we also require that the 12 second time intervals before the flare start time and after the flare end time be "clean", so that we can use those times to calculate the background level. This process resulted in a sample of 10,661 flares. Spectra were fit for each flare for the same time range that was used to find the [http://sprg.ssl.berkeley.edu/~jimm/hessi/flare_list_20061004 flare position]. This time interval is usually an interval of 2 minutes or less (depending on flare duration) at the flare peak in the 6 to 12 keV range. The spectra were assumed to be isothermal, and fit for the energy range from 6 to 20 keV, using an energy resolution of 1/3 keV. <br />
<br />
Not all of the flares were fit successfully for all detectors; flares which did not have count rates that were more than 3 sigma above the background level in more than 2 channels were discarded; flares for which the calculated EM was less than 10<sup>43</sup> cm<sup>-3</sup> were discarded; flares with background-subtracted count rates totaled over the 6 to 20 keV range which were less than zero were discarded. These tests were applied for each of the detectors 3, 4, 5, 6, 8, and 9. Spectra from detectors 2 and 7 have been ignored, since those detectors have reduced energy resolution in the 6 to 20 keV range.<br />
<br />
The full sample contains 7,740 flares, from the start of the mission until 1 November 2009.<br />
<br />
== Results for relative sensitivity ==<br />
<br />
[[Image:Hessi_dets_test.png|300px|thumb|center|'''Figure 1''':This is the relative detector efficiency for front detectors 3, 4, 5, 6, 8, and 9 ]]<br />
<br />
Figure 1 shows 60-day averages of the ratio between the EM of each of detectors 3, 4, 5, 6, 8, and 9 and detector 1 for the full time range. From the plot we see that relative sensitivity was closely the same for all detectors until 2006. Early in 2006, the relative sensitivity of detector 3 began to drop. Later in 2006 the detector 5 sensitivity rolls over. Detectors 4, 6, 8, and 9 lose sensitivity during 2007. The first data gap on the plot is for the detector [http://sprg.ssl.berkeley.edu/~tohban/nuggets/?page=article&article_id=69 anneal] which took place in November 2007. Post-anneal, the detectors recover sensitivity. The second data gap on the plot is just a quiet time. After the second gap, we can see detectors losing sensitivity again.<br />
<br />
Note that the relative sensitivity of all detectors seems to be greater than unity for the early part of the mission. This may be a systematic effect resulting from holding the temperature fixed at the level found using detector 1 for each of the detectors, i.e. that something about detector 1 leads to a biased temperature value. This possibility is being investigated. <br />
<br />
== What about absolute sensitivity? ==<br />
<br />
For each flare we also measured [http://www.swpc.noaa.gov/today.html#xray GOES] flux averaged for the spectral interval. We expect that the GOES detectors are less susceptible to decay then the RHESSI detectors, which suffer from radiation damage. For most of the flares GOES 10 data were used, but there are gaps in the GOES 10 coverage, so there are flares with GOES 11 and 12 data included. (It turns out that even if non-GOES 10 flares are discarded, we come to the same overall conclusion.) Figure 2 shows 60 day averages of the ratio of RHESSI detector 1 emission measure to GOES fluxes for the full sample, which should show the variation of the relative sensitivity of detector 1 to GOES. Since we see that the other detectors lose sensitivity, we expect that detector 1 should lose some sensitivity, but it seems to gain sensitivity relative to the GOES detectors. How is this possible?<br />
<br />
<br />
{| border="0" align="center"<br />
|+<br />
|-<br />
| valign="top"|<br />
[[Image:Hessi_dets_test1.png|300px|thumb|right|'''Figure 2''':60 day averages of the ratio of RHESSI detector 1 EM to GOES fluxes for the full sample.]]<br />
| valign="top"|<br />
[[Image:Hessi_dets_test0.png|300px|thumb|left|'''Figure 3''':GOES Hi channel flux versus RHESSI detector 1 emission measure for the full sample.]]<br />
|-<br />
|}<br />
<br />
It turns out that there are some effects that are not being accounted for. Figure 3 shows the GOES high-energy, short-wavelength channel flux versus RHESSI detector 1 EM for the full sample. We expect correlation between these two quantities, considering that both instruments are observing X-rays, but with different energy ranges; the GOES energy response is weighted towards energies less than 6 keV, while the RHESSI energy response is small at these low energies and increased at higher energy, where GOES is insensitive. Also the correlation depends on the shape of the observed photon spectrum, which here is assumed to be due to an isothermal source. We are fairly sure that<br />
flares are not isothermal (Ref. 1), and this non-isothermality may be responsible for some of the scatter on the plot.<br />
<br />
There is correlation for relatively large flares, but there is quite a bit of scatter, and for small events, there is no correlation at all. There is a lower limit for GOES, and there are line-like structures in the GOES fluxes which show the effects of digitization in the GOES data. We expect that some of the scatter is due to lack of background subtraction for the GOES data. (Background subtraction for GOES will be included in the next iteration of this project.) This will be a smaller effect for larger flares. Also we may hope to reduce scatter by comparing GOES flux to RHESSI count rate, rather than EM since there is one less level of data processing involved. So for the comparison shown in Figure 4 we restrict the number of flares by only taking relatively large flares, with EM > 10<sup>45</sup> cm<sup>-3</sup>, which is the value below which the correlation between flux and EM disappears. This leaves 2,580 flares. Also we will compare GOES flux with RHESSI count rate, instead of EM. <br />
<br />
Once we have made these changes, we get a good correlation, as shown in Figure 4. The red dashed line shows a linear least-squares fit to the data with a slope of 0.55. There is still scatter; we expect that this will decrease when the GOES background subtraction is included. If we plot the 60-day averages of the ratio of RHESSI detector 1 counts to GOES flux for this flare sample, we obtain the results shown in Figure 5. Note that the error bars in both Figures 2 and 5 are calculated using the standard deviation of the mean of the ratios in each time bin, which is the standard deviation divided by the square root of the number of flares per time bin. There are more flares per bin during the early part of the mission (typically 100 flares/bin) than later (often less than 10 flares/bin); thus the error bars are smaller earlier in the mission. Also there are now three data gaps, one for annealing, and two due to lack of activity. <br />
<br />
{| border="0" align="center"<br />
|+<br />
|-<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct04.png|300px|thumb|left|'''Figure 4''':GOES Hi channel flux versus RHESSI detector 1 6 to 20 keV count rate for the flares with EM greater than 10<sup>45</sup> cm<sup>-3</sup>.]]<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct14.png|300px|thumb|right|'''Figure 5''':60 day averages of the ratio of RHESSI 6 to 20 keV counts to GOES fluxes for the small sample of flares with EM greater than 10<sup>45</sup> cm<sup>-3</sup>.]]<br />
|-<br />
|}<br />
<br />
In Figure 5 the variation of RHESSI counts with respect to GOES flux is much less than the variation shown in Figure 1. The level now looks relatively constant. There is no evidence from this plot that RHESSI detector 1 has lost or gained sensitivity with respect to GOES. It is also probable here that subtracting the GOES background for each flare will affect these results. This will be addressed soon.<br />
<br />
== How does the sample change affect the original plot? ==<br />
<br />
[[Image:Hessi_dets_test_ct4.png|300px|thumb|center|'''Figure 6''':This is the relative detector efficiency for front detectors 3, 4, 5, 6, 8, and 9, but using count rates and only the small sample of 2,580 flares with EM greater than 10<sup>45</sup> cm<sup>-3</sup>.]]<br />
<br />
Using only relatively large flares in the sample, and using count rate rather than EM for the comparison has a small effect on the relative sensitivity plot, as shown in Figure 6. Except for the fact the the detector 3 curve is less than one at the start of the mission, and the odd large value for detector 8 near the end of 2008, not much has changed.<br />
<br />
== Conclusions ==<br />
<br />
Detectors 3, 5, 6, 8, and 9 lose sensitivity relative to detector 1, especially in 2007. Some of that sensitivity is regained by annealing, but it is currently decreasing again.<br />
We will be able to anneal the RHESSI detectors again, and expect as before to restore at least part of their original performance, but we want to do this at the optimum time. <br />
This calibration exercise will help us to understand when that time will come. <br />
<br />
There is no obvious loss of sensitivity in detector 1, relative to GOES. It seems odd that the other detectors can lose sensitivity, so why not detector 1 as well? The next step for this work is to try to subtract pre-flare background for the GOES flare results. Watch this space for updates.<br />
<br />
== Update: add GOES background subtraction <font color=#0000f0> (3 January 2010) </font> ==<br />
<br />
Here are what the results look like when you do subtract GOES background. As a first step, we subtracted pre-flare background values from the GOES fluxes before comparing to RHESSI count rates. For each flare the pre-flare background time interval used is the same as the pre-flare interval found for the RHESSI spectra. (Note that this is slightly different than the background subtraction process for RHESSI, which included both pre- and post-flare background levels and interpolated. Since GOES flares tend to last longer than RHESSI flares, we are concerned that using both the pre- and post flare intervals calculated from RHESSI data may give inaccurate results.) It turns out that only 2055 flares had excess background-subtracted GOES flux. The results for this sample are shown in Figures 7 and 8. In Figure 7, the scatter in the plot of GOES flux versus RHESSI counts is much less than that in Figure 4, as expected. The red dashed line is a least-squares fit to the data, with a slope of 0.86. <br />
<br />
{| border="0" align="center"<br />
|+<br />
|-<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct0b.png|300px|thumb|left|'''Figure 7''':Background-subtracted GOES Hi channel flux versus RHESSI detector 1 6 to 20 keV count rate.]]<br />
| valign="top"|<br />
[[Image:Hessi_dets_test_ct1b.png|300px|thumb|right|'''Figure 8''':60 day averages of the ratio of RHESSI 6 to 20 keV counts to background-subtracted GOES fluxes.]]<br />
|-<br />
|}<br />
<br />
In Figure 8, which shows the variation of the ratio of RHESSI counts to background-subtracted GOES flux, the relative sensitivity finally shows a decrease. Relative to GOES Hi, RHESSI detector 1 loses about half of its sensitivity by the annealing in November 2007. This is shown a little more clearly in Figure 9, which normalizes the ratio to be 1 at the start of the mission. As in the other figures, sensitivity increases after annealing (first datagap), has been decreasing since, and is now at a level of approximately half of the original sensitivity. <br />
<br />
[[Image:Hessi_dets_test_hi_ct1b.png|300px|thumb|center|'''Figure 9''':This is the relative detector 1 efficiency normalized to be 1.0 for the first 60 days of the mission, using background-subtracted GOES fluxes for the comparison.]]<br />
<br />
== New Conclusion ==<br />
<br />
There is now an obvious loss of sensitivity in detector 1, relative to GOES, which we find by including pre-flare background subtraction for the GOES flare results.<br />
<br />
== References ==<br />
1. http://adsabs.harvard.edu/abs/1999ApJ...514..472M The Solar Flare Soft X-Ray Differential Emission Measure and the Neupert Effect at Different Temperatures by J.McTiernan, G. Fisher and P. Li<br />
<br />
[[Has observation by:: RHESSI| ]]<br />
[[Has article subject:: RHESSI]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Understanding_the_co-spatial_return_current_in_solar_flaresUnderstanding the co-spatial return current in solar flares2018-09-18T02:17:09Z<p>Schriste: /* The return-current collisional thick-target model */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title = Understanding the co-spatial return current in solar flares <br />
|number = 330<br />
|first_author = Meriem Alaoui <br />
|second_author = Gordon Holman<br />
|publish_date = 6 August 2018<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::331]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::329]]}}<br />
}}<br />
<br />
== Introduction ==<br />
<br />
In the standard description of flare hard X-ray emission, an <br />
unspecified mechanism accelerates large numbers of electrons in the<br />
corona.<br />
These then penetrate into the lower atmosphere, creating "footpoint" <br />
sources of emission across the spectrum.<br />
These <br />
[https://en.wikipedia.org/wiki/Particle_beam beams] require <br />
return currents to balance the charge distribution.<br />
Co-spatial return currents have been proposed for this purpose.<br />
The return current locally neutralizes the charge build-up<br />
and cancels the magnetic field induced by the beam of accelerated<br />
electrons. <br />
Return currents solve the so-called number problem and<br />
the associated current stability problem. <br />
In this Nugget we describe the constraints on the flare-nonthermal<br />
electron distribution and its associated return current using RHESSI<br />
x-ray imaging-spectroscopy observations of many flares (Ref. [1]) .<br />
<br />
[[File:330f1.png|400px|thumb|center|Figure 1: <br />
Illustration of the return-current thick-target model. <br />
Deceleration of the beam electrons by the return-current electric field <br />
dominate energy losses in the corona, and Coulomb collisions dominate <br />
at the footpoints.<br />
]]<br />
<br />
== The return-current collisional thick-target model ==<br />
<br />
The flares studied in this work exhibit strong spectral flattenings <br />
of the x-ray spectrum at lower energies (a few deka-keV) during the<br />
HXR peak times. <br />
The magnitude of the flattening cannot be explained<br />
by mechanisms such as non-uniform ionization, isotropic Compton<br />
back-scattering, or pulse pile-up alone (see details in Ref. [1]).<br />
<br />
We consider a model (Figure 1) based on coronal magnetic loops, with <br />
the following assumptions:<br />
<br />
# Electrons are accelerated at the looptop with a single power-law spectrum and a sharp low-energy cutoff (E<sub>c).<br />
# The electrons lose some of their energy by return-current losses along the path toward the thick target.<br />
# They lose all of their remaining energy by Coulomb collisions. <br />
# We have a 1D model with electrons streaming along the loop and a co-spatial return current in the opposite direction.<br />
# The return current has had time to reach a steady state within the integration time of the spectra.<br />
# All the X-ray spectral flattening is due to return-current losses. This provides the upper limit of the return-current potential drop from the acceleration region to the footpoints. <br />
# Electrons are thermalized and lost from the beam when their energy reaches the lower limit of the cutoff energy.<br />
<br />
== Are the beam and the return current stable? ==<br />
<br />
We have applied this model to a large number of flares, <br />
including some "early-impulsive" events (Ref. [2]) for which the non-thermal<br />
hard X-ray emission can be detected at lower energies, and others with clear <br />
flattening at low energies (the "spectral break") was clear.<br />
In each case we have evaluated the criterion for stability against the<br />
generation of<br />
[https://en.wikipedia.org/wiki/Electrostatic_ion_cyclotron_wave electrostatic ion-cyclotron waves]<br />
as well as ion acoustic waves and the Buneman instability, for which the threshold is about 2v<sub>e</sub\>.<br />
For each case we can evaluate the drift velocity, assuming the beam<br />
current to be exactly balanced by the return current, and compare that<br />
with the known instability threshold (Figure 2).<br />
The comparison is made for separate estimates of upper and lower limits<br />
for the cutoff energy of the electron spectrum.<br />
<br />
[[File:330f2.png|600px|thumb|center|Figure 2:<br />
The ratio of the return-current drift velocity to the ion acoustic<br />
speed as a function of the electron to ion temperature ratio for<br />
the upper and lower limits estimated for the cutoff energy E<sub>c</sub>.<br />
The X-axis shows different assumptions about the ratio of electron to ion<br />
temperature in the source.<br />
In most cases, the return current drift velocity is too low to<br />
generate electrostatic ion cyclotron and ion acoustic instabilities<br />
]]<br />
<br />
The resistivity is enhanced in most cases, if Ohm’s law is valid (figure 3). <br />
Current-driven instabilities are not the cause of that enhancement (figure 2). <br />
It is possible that the resistivity does not need to be enhanced, if part of the return current is carried by freely accelerated electrons. <br />
Ohm’s law is not valid then for calculating the resistivity, and the return-current electric field is on the order of the Dreicer field, which means runaway electrons cannot be neglected.<br />
<br />
[[File:330f3.png|600px|thumb|center|Figure 3:<br />
The derived resistivity is enhanced by more than 2 orders of magnitude in most cases. <br />
The figure shows scatter plots of the derived resistivity in the upper (left) and lower (right) limits of the low-energy cutoff vs. the classical Spitzer resistivity. <br />
The error bars are the 67% confidence intervals. <br />
The two dashed lines on the left (right) panel are where &eta;<sub>spitzer</sub> = &eta; (E<sub>c</sub> max) and &eta;(E<sub>c</sub> max)= 10<sup>4</sup> x &eta;<sub>spitzer</sub> (&eta;(E<sub>c</sub> max)= 10<sup>7</sup> x &eta;<sub>spitzer</sub>). <br />
The green horizontal lines in both panels show the Spitzer resistivity of a 2 MK plasma. <br />
Since most of the data points are higher than the green line, it means that even if the beam is streaming along a cooler loop than the one observed with RHESSI, which is most likely from previously heated plasma, the resistivity values are still enhanced in most cases, as compared to Spitzer values.<br />
]]<br />
<br />
== Conclusions ==<br />
<br />
A main conclusion of the statistical analysis is that if the<br />
flattening is caused by return-current energy losses, then return<br />
currents are unlikely to be carried by only the bulk of the thermal<br />
distribution (electrons with velocities near the thermal speed, v<sub>e</sub>), <br />
but if the return current is carried by the bulk thermal electrons, <br />
the resistivity must be<br />
enhanced by a process different from the usual current-driven<br />
instabilities, namely, ion acoustic, electrostatic ion cyclotron,<br />
and Buneman instabilities. <br />
This was inferred from the result that<br />
the drift velocity of the return current was found to be lower than<br />
the threshold for generating these instabilities in all cases using<br />
the upper limit of the low-energy cutoff E<sub>c</sub>, <br />
and in 506 out 528 cases using the lower limit of E<sub>c</sub>.<br />
<br />
Based on the results that these current-driven instabilities cannot<br />
explain the enhanced (anomalous) resistivity values, and that the<br />
return-current electric field is on the order of the Dreicer field,<br />
when the resistivity is classical, we will examine in an upcoming<br />
paper a scenario where the return-current carrying electrons are<br />
in the runaway regime, and whether the need for anomalous resistivity<br />
can be abandoned.<br />
<br />
== References ==<br />
<br />
[1] [http://adsabs.harvard.edu/abs/2017ApJ...851...78A "Understanding Breaks in Flare X-Ray Spectra: Evaluation of a Cospatial Collisional Return-current Model"] <br />
<br />
[2] [http://adsabs.harvard.edu/abs/2007ApJ...670..862S "Nonthermal X-Ray Spectral Flattening toward Low Energies in Early Impulsive Flares"]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Help:For_AuthorsHelp:For Authors2018-09-18T02:15:26Z<p>Schriste: </p>
<hr />
<div>This page describes the steps necessary to create and edit your own article. This information is generally valid for creating any kind of article on this wiki but may be particularly useful to new Nugget authors.<br />
<br />
== Creating an Article ==<br />
All pages on this wiki are articles including the RHESSI science nuggets. In order to create a new article a title must first be chosen. The title of an article cannot be changed so choose your article title carefully! It is possible to move the content from one article to another with a different title but this is not a recommended practice. There are two main ways to create an article.<br />
<br />
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Also please be aware that many of the audience are not native English speakers, so phrases like "X must have been a gutsy operator to have diagonalized those macrospicules" would not do so well.<br />
Writing as though for a newspaper, rather than as for a boring archival journal, would be best - there is no need really to have complete literature <br />
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==Meta data==<br />
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In order to provide better tools to organize and search for nugget topics, we encourage authors to use the following meta data tags in their articles. <br />
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* To note that an article discusses a particular event use the following <pre>[[Has event date:: November 09, 2002 13:12]]</pre><br />
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The following subjects are available.<br />
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== Source Code ==<br />
If you'd like to add actual source code in your article please use the following syntax.<br />
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<br />
using System;<br />
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public static int Main(String[] args)<br />
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Console.WriteLine("Hello, World!");<br />
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</source></div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Help:For_AuthorsHelp:For Authors2018-09-18T02:07:05Z<p>Schriste: /* Source Code */</p>
<hr />
<div>This page describes the steps necessary to create and edit your own article. This information is generally valid for creating any kind of article on this wiki but may be particularly useful to new Nugget authors.<br />
<br />
== Creating an Article ==<br />
All pages on this wiki are articles including the RHESSI science nuggets. In order to create a new article a title must first be chosen. The title of an article cannot be changed so choose your article title carefully! It is possible to move the content from one article to another with a different title but this is not a recommended practice. There are two main ways to create an article.<br />
<br />
*Simply search for the title of the article you wish to create using the search bar on the left. The wiki should confirm that no other articles have this title by not finding it and will give a choice to create this article. Click the link (at "you searched for") and your new article will open in edit mode.<br />
<br />
*Articles are organized through categories. An article is added to a category by adding meta-data into the article text. The next method for creating an article is to navigate to the category page and use the Form to create a new article in that category. The RHESSI Science Nugget category page can be found [[:Category:Nugget|here]]. The category meta-data will then automatically be included in the article text.<br />
<br />
== Writing your article ==<br />
<br />
Writing a wiki article is very straightforward. Every article has an "edit" tab above it. Clicking that will open in the article in an editable window. Many formatting shortcuts can be found above the editing window. Since each article on this wiki can be edited, it is easy to learn by example. For more basic help with editing read the following [[Help:Editing|article]].<br />
<br />
=== Adding Figures ===<br />
In order to add a figure it is first necessary to upload the image to the wiki. This can be done by using the link on the left titled "Upload File". Please use a descriptive name for the file and add a description of its contents. Once uploaded the file can be added to articles using the following command<br />
'''<nowiki>[[Image:POF.jpg|200px|thumb|right|'''Figure 1''': This is the figure caption.]]</nowiki>'''<br />
whose results can be seen to the right. Most of these parameters should be self explanatory. The "thumb" parameters tells the wiki that this is a figure and therefore displays the figure caption. The float parameters (in this case set to "right") tells the wiki where to float the image. Possibilities include "left", "center", or "right". <br />
<br />
[[Image:POF.jpg|250px|thumb|right|'''Figure 1''': This is the figure caption. Note that the figure has properlyy appeared on the right, but that you can't really control how the text will wrap around it - that depends on the browser.]]<br />
<br />
If you'd like to experiment with more advanced methods read this [http://en.wikipedia.org/wiki/Image_tutorial article] on Wikipedia's website.<br />
<br />
=== Style and content ===<br />
<br />
We try to write the Nuggets so that a technically literate reader won't be baffled.<br />
That means that the text should appear in plain English, and that jargon is a no-no. <br />
Many scientists don't realize they are writing gibberish so please be careful about this. <br />
On the plus side, you can freely use cgs units if you wish.<br />
Also please be aware that many of the audience are not native English speakers, so phrases like "X must have been a gutsy operator to have diagonalized those macrospicules" would not do so well.<br />
Writing as though for a newspaper, rather than as for a boring archival journal, would be best - there is no need really to have complete literature <br />
citation, since those ''really'' knowledgeable about the field will certainly know where to go (ADS).<br />
<br />
== Source Code ==<br />
If you'd like to add actual source code in your article please use the following syntax.<br />
<br />
<source lang="csharp"><br />
// Hello World in Microsoft C# ("C-Sharp").<br />
<br />
using System;<br />
<br />
class HelloWorld<br />
{<br />
public static int Main(String[] args)<br />
{<br />
Console.WriteLine("Hello, World!");<br />
return 0;<br />
}<br />
}<br />
</source></div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/RHESSI_Science_NuggetsRHESSI Science Nuggets2018-09-18T02:05:55Z<p>Schriste: </p>
<hr />
<div>Welcome to the [[RHESSI Science Nuggets]]: science notes from [[RHESSI]]. The following is a time-ordered list of the latest Nuggets added to the wiki. An [[:Category:Nugget|alphabetical list of wiki Nuggets]] is also available as well as [[:Category:RHESSI Nugget List|yearly lists]]. We welcome volunteer authors - please see our page of [[Help:For_Authors| help for authors]].<br />
<br />
{{Nugget Badge<br />
|title = New Views of Global Solar Magnetic Field Evolution Over Four Solar Cycles<br />
|number = 331<br />
|first_author = David Webb<br />
|second_author = <br />
|publish_date = 27 August 2018<br />
|description = A digital archive of Pat McIntosh's 44 years of solar synoptic observations <br />
|image=Icon331.png}}<br />
<br />
{{Nugget Badge<br />
|title = Understanding the co-spatial return current in solar flares<br />
|number = 330<br />
|first_author = Meriem Alaoui<br />
|second_author = and Gordon Holman<br />
|publish_date = 6 August 2018<br />
|description = Completing the circuit in a thick-target model <br />
|image=Icon330.png}}<br />
<br />
{{Nugget Badge<br />
|title = 3D Magnetic Reconnection at a Coronal Null Point<br />
|number = 329<br />
|first_author = Shane Maloney,<br />
|second_author = Aidan O'Flannagain and Peter Gallagher<br />
|publish_date = 30 July 2018<br />
|description = Large-scale reconnection involved in Type I radio noise storm <br />
|image=Icon329.png}}<br />
<br />
{{Nugget Badge<br />
|title = The true dawn of multimessenger astronomy<br />
|number = 328<br />
|first_author = Hugh Hudson<br />
|second_author = <br />
|publish_date = 23 July 2018<br />
|description = Ever since the Carrington flare <br />
|image=Icon328.png}}<br />
<br />
{{Nugget Badge<br />
|title = Microwave Imaging Spectroscopy of Flares is Here<br />
|number = 327<br />
|first_author = Dale E. Gary,<br />
|second_author = EOVSA and RHESSI Teams<br />
|publish_date = 16 July 2018<br />
|description = Microwave imaging spectroscopy takes a giant leap forward with SOL2017-09-10 <br />
|image=Icon327.png}}<br />
<br />
{{Nugget Badge<br />
|title = Coronal nanoflares powered by footpoint reconnection<br />
|number = 326<br />
|first_author = Pradeep Chitta,<br />
|second_author = Hardi Peter, and Sami Solanki<br />
|publish_date = 9 July 2018<br />
|description = Coronal nanoflares in active region cores can be powered by the magnetic reconnection in the lower solar atmosphere <br />
|image=Icon326.png}}<br />
<br />
{{Nugget Badge<br />
|title = A remarkable, but confused, coronal hard X-ray source<br />
|number = 325<br />
|first_author = Alexandra Lysenko,<br />
|second_author = Larisa Kashapova and Hugh Hudson<br />
|publish_date = 25 June 2018<br />
|description = A remarkable flare in 1999 adds to our short list of extended coronal hard X-ray/microwave sources <br />
|image=Icon325.png}}<br />
<br />
{{Nugget Badge<br />
|title = Understanding HMI pseudocontinuum in white-light flares<br />
|number = 324<br />
|first_author = Michal &Scaron;vanda<br />
|second_author = et al.<br />
|publish_date = 28 May 2018<br />
|description = The HMI pseudocontinuum (Ic) is ill-calibrated in regions with strong fields, i.e. for white-light flares <br />
|image=Icon324.png}}<br />
<br />
{{Nugget Badge<br />
|title = To beam or not to beam - that is (still) the question<br />
|number = 323<br />
|first_author = Paulo Sim&otilde;es<br />
|second_author = and Hugh Hudson<br />
|publish_date = 14 May 2018<br />
|description = Descriptions of the lower solar atmosphere of flares <i>ca.</i> Cycle 21 sound surprisingly current <br />
|image=Icon323.png}}<br />
<br />
{{Nugget Badge<br />
|title = Observation of Cosmic Ray Spallation Events from SoHO<br />
|number = 322<br />
|first_author = Serge Koutchmy<br />
|second_author = and Ehsan Tavabi<br />
|publish_date = 7 May 2018<br />
|description = LASCO's images capture high-energy nuclear interactions from cosmic-ray hits <br />
|image=Icon322.png}}<br />
<br />
{{Nugget Badge<br />
|title = A Sunspot from Cycle 25 for sure<br />
|number = 321<br />
|first_author = Tomek Mrozek<br />
|second_author = and Hugh Hudson<br />
|publish_date = 10 April 2018<br />
|description = YES! Cycle 25 is here! <br />
|image=Icon321.png}}<br />
<br />
{{Nugget Badge<br />
|title = Blue-wing enhancement of the Mg II h and k lines in a flare<br />
|number = 320<br />
|first_author = Akiko TEI<br />
|second_author =<br />
|publish_date = 9 April 2018<br />
|description = Flare loops involve a cool upflow preceding the hot evaporation flow <br />
|image=Icon320.png}}<br />
<br />
{{Nugget Badge<br />
|title = NuSTAR detects X-ray flares in the quiet Sun<br />
|number = 319<br />
|first_author = Matej Kuhar<br />
|second_author = and S&auml;m Krucker<br />
|publish_date = 26 March 2018<br />
|description = Quiet-Sun flares may not be powerful, but they look a lot like ordinary flares<br />
|image=Icon319.png}}<br />
<br />
{{Nugget Badge<br />
|title = Homologous CME/flares from AR 12371<br />
|number = 318<br />
|first_author = Panditi Vemareddy<br />
|second_author = and Pascal Demoul&iacute;n<br />
|publish_date = 19 March 2018<br />
|description = An excellent set of homologous flare/CMEs analyzed and explained<br />
|image=Icon318.png}}<br />
<br />
{{Nugget Badge<br />
|title = Non-Maxwellian Diagnostics from SDO/EVE Spectra of an X-class Flare<br />
|number = 317<br />
|first_author = Elena Dzif&#x10d;&aacute;kov&aacute;<br />
|second_author = and Jaroslav Dud&iacute;k<br />
|publish_date = 16 February 2018<br />
|description = Ratios of high-excitation ions can readily detect &kappa;-distributions in flare plasmas<br />
|image=Icon317.png}}<br />
<br />
{{Nugget Badge<br />
|title = Joint MinXSS and RHESSI Flare X-ray Spectra between 1 and 15 keV<br />
|number = 316<br />
|first_author = Chris Moore, Brian Dennis and the MinXSS Science Team<br />
|publish_date = 5 February 2018<br />
|description = MinXSS adds systematic views of flare soft X-ray spectra to RHESSI imagery<br />
|image=Icon316.png}}<br />
<br />
{{Nugget Badge<br />
|title = Parameterized Flare Models with Chromospheric Compressions<br />
|number = 315<br />
|first_author = Adam Kowalski & Joel Allred<br />
|publish_date = 17 January 2018<br />
|description = A new approach to modeling the lower flare atmosphere<br />
|image=FlareModelsKowalskiAllred.png}}<br />
<br />
{{Nugget Badge<br />
|title = A Curious Sunspot Group in 2018<br />
|number = 314<br />
|first_author = Hugh Hudson<br />
|publish_date = 14 January 2018<br />
|description = The first new sunspot group of 2018 emerged at the wrong latitude<br />
|image = Icon314.png}}<br />
<br />
{{Nugget Badge<br />
|title = Tecumseh's Eclipse and Astrophysics<br />
|number = 313<br />
|first_author = Hugh Hudson<br />
|publish_date = 25 December 2017<br />
|description = The solar corona was first recognized as such, and named, in an eclipse of 1806<br />
|image = Icon313.png}}<br />
<br />
{{Nugget Badge<br />
|title = Hunting for Hidden Tiny Flares<br />
|number = 312<br />
|first_author = Shin-nosuke ISHIKAWA<br />
|publish_date = 27 November 2017<br />
|description = FOXSI-2 says that episodic energy releases are still viable as a part of the coronal heating problem.<br />
|image = Icon312.png}}<br />
<br />
{{Nugget Badge<br />
|title = Unusual Type III Burst Dynamics Produced by Diverging Magnetic Fields<br />
|number = 311<br />
|first_author = Patrick McCauley<br />
|publish_date = 20 November 2017<br />
|description = Unusual type III bursts follow coronal separatrix structures.<br />
|image = Icon311.png}}<br />
<br />
{{Nugget Badge<br />
|title = Valderrama in the 21st Century<br />
|number = 310<br />
|first_author = Hugh Hudson<br />
|publish_date = 31 October 2017<br />
|description = A newly-described white-light flare from the 19th century!..<br />
|image = Icon310.png}}<br />
<br />
{{Nugget Badge<br />
|title = Electron Scattering in the Flaring Corona<br />
|number = 309<br />
|first_author = Sophie Musset<br />
|publish_date = 24 October 2017<br />
|description = Diffusive transport may contribute to the trapping of electrons in coronal X-ray sources <br />
|image = Icon309.png}}<br />
<br />
{{Nugget Badge<br />
|title = The Power of Turbulence<br />
|number = 308<br />
|first_author = Nic Bian<br />
|publish_date = 25 September 2017<br />
|description = Turbulent energy content may underlie flare energy transfer, magnetic reconnection, and particle acceleration <br />
|image = Icon308.png}}<br />
<br />
{{Nugget Badge<br />
|title = The Kelvin Force and Loop-Top Concentration<br />
|number = 307<br />
|first_author = Kiyoto SHIBASAKI<br />
|publish_date = 18 September 2017<br />
|description = New physics can explain the perplexing overpressure at the flare looptop regions<br />
|image = Icon307.png}}<br />
<br />
{{Nugget Badge<br />
|title = The Last Best Flare of Cycle 24?<br />
|number = 306<br />
|first_author = S&auml;m Krucker<br />
|second_author = and Hugh Hudson<br />
|publish_date = 11 September 2017<br />
|description = Right on schedule, Cycle 24 has produced a great flare (with a GLE)<br />
|image = Icon306.png}}<br />
<br />
{{Nugget Badge<br />
|title = Electric Current Neutralization and Solar Eruption in Active Regions<br />
|number = 305<br />
|first_author = Yang LIU<br />
|second_author = <br />
|publish_date = 28 August 2017<br />
|description = Active current systems in the solar corona don't have return currents<br />
|image = Icon305.png}}<br />
<br />
{{Nugget Badge<br />
|title = RHESSI and the Megamovie<br />
|number = 304<br />
|first_author = Hugh Hudson, Laura Peticolas,<br />
|second_author = and Juan Carlos Mart&iacute;nez Oliveros<br />
|publish_date = 31 July 2017<br />
|description = A wholly new way to view a solar eclipse, and to do solar astrometry<br />
|image = Icon304.png}}<br />
<br />
{{Nugget Badge<br />
|title = Bastille Day 2017<br />
|number = 303<br />
|first_author = Hugh Hudson<br />
|second_author = and S&auml;m Krucker<br />
|publish_date = 24 July 2017<br />
|description = Interesting flares really do happen on Bastille Day...<br />
|image = Icon303.png}}<br />
<br />
{{Nugget Badge<br />
|title = The Solar X-ray Limb III<br />
|number = 302<br />
|first_author = Marina Battaglia<br />
|second_author = and Gordon Hurford<br />
|publish_date = 12 June 2017<br />
|description = RHESSI succeeds with a wholly new way to measure the solar diameter<br />
|image = Icon302.png}}<br />
<br />
{{Nugget Badge<br />
|title = Double Coronal X-ray and Microwave Sources Associated With A Magnetic Breakout Solar Eruption<br />
|number = 301<br />
|first_author = Yao CHEN<br />
|second_author =<br />
|publish_date = 29 May 2017<br />
|description = A different explanation of the double coronal hard X-ray sources<br />
|image = Icon301.png}}<br />
<br />
{{Nugget Badge<br />
|title = A Lasso Model for Solar Gamma-ray Events<br />
|number = 300<br />
|first_author = Hugh Hudson<br />
|second_author =<br />
|publish_date = 15 May 2017<br />
|description = A toy model hoping to explain the SEP/LAT relationship<br />
|image = Icon300.png}}<br />
<br />
[[RHESSI Science Nuggets 200 to 299|Next Nuggets]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Birth_of_a_dense_flaring_loopBirth of a dense flaring loop2018-09-17T19:00:11Z<p>Schriste: /* Conclusions */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Birth of a dense flaring loop<br />
|number=23<br />
|first_author=Laura Bone<br />
|publish_date=27 January 2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::24]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::22]]}}<br />
}}<br />
<br />
<br />
==Introduction==<br />
<br />
In the classical "thick target" model of a solar flare, high-energy electrons are accelerated at some reconnection site in the corona. The electrons stream down the magnetic field lines to the chromosphere where they are stopped by the dense plasma and produce hard X-ray emission by collisional bremsstrahlung. The plasma at the footpoints then heats up, and "evaporates" into the coronal magnetic loop, producing thermal soft X-ray emission.<br />
<br />
From the early 1970's hard X-ray emission has occasionally been seen in the corona as well, initially in occulted events (flares with the footpoints behind the limb). These first events were interpreted to be thermal bremsstrahlung from hot coronal plasma, with significant thick target emission only at the footpoints. The first observation of an impulsive, and by implication non-thermal coronal source, was the famous <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/An_Alternative_View_of_the_Masuda_Flare Masuda flare]. <br />
Later Yohkoh studies of X-ray flares demonstrated that coronal sources exhibit both gradual and impulsive hard X-ray emission.<br />
<br />
Since its launch in 2002 RHESSI, has observed a number of both moving and stationary coronal sources. The frequency of these sources and the diversity of the physics they show us has been one of the major developments of the RHESSI era. Theoretically one of the simplest explanations for coronal sources, particularly in flares with little or no footpoint emission, is an enhanced coronal density, but then we must explain how the density came to be so large in the first place. We must also ask whether these dense regions can be observed in other ways, e.g. via the techniques of solar radio astronomy.<br />
<br />
== SOL2002-04-14 ==<br />
<br />
We have studied a flare in which RHESSI observes emission up to 50 keV, definitely in the "hard" or "non-thermal" energy range. This unocculted M3.7 flare, on the [[Has event date:: April 14, 2002 00:05]]<br />
, showed hard X-ray emission primarily from the corona, with little or no footpoint emission during most of the event (see Figure 1). This was interpreted in terms of high coronal column densities: electrons up to ~50 keV actually stopped high up in the corona, in effect producing a coronal thick target.<br />
<br />
[[File:23f1.png|600px|thumb|center|Figure 1: RHESSI images at 6-12keV with overlayed contours at 25-50 keV. Footpoint emission is only visible in the higher energy band for a brief period during the flare maximum, around 00:10 UT. During the rest of the flare the X-ray emission is confined primarily to the loop top (image courtesy of Astrid Veronig).]]<br />
<br />
At the onset of this event emission measures derived from the GOES soft X-ray data suggested that the column density was already rather high, ~1020cm-2, implying that the loop was already filled with plasma. For this flare complementary multi-wavelength observations provided us with information allowing us to understand the active region evolution and sequence of events in which the dense loop was born. Two flares occurred in this active region in close temporal and spatial proximity and this sequence of flares leads us to propose a series of reconnections between multiple loops in a 3-dimensional field geometry.<br />
<br />
==Model of the loop interaction==<br />
<br />
The loop in which the M3.7 flare occurs can be seen to brighten gradually both in EIT 195 Å and 17 GHz (radio microwave band) just after the impulsive phase of the earlier flare. Combined with the already bright, polarised radio emission from the southern footpoint, this is evidence that fast electrons may already have been present in the flaring loop (see Figure 2).<br />
<br />
[[File:23f2.png|600px|thumb|center|Figure 2: The evolution of the microwave emission (grayscale images) over the course of the flare; contours show radio polarization. The first panel shows the latter half of the earlier flare. The second image is from between the flares. The third panel shows the whole loop in which the M3.7 flare occurs beginning to brighten and the fourth image shows the whole loop at the peak of the flare (image courtesy of Stephen White).]]<br />
<br />
The earlier flare occurs at the interface of two active regions through reconnection it forms a small dense loop between the two active regions and an area of overlying field. The small loop, observed to be initially rather dense, becomes unstable; this leads to further particle acceleration and density increase and the rise of the loop which interacts with the overlying field (see Figure 3).<br />
<br />
[[File:23f3.png|600px|thumb|center|Figure 3: The sequences of reconnection occurring between the magnetic loops in the two flares. The initial reconnection forms the overlying loops and the smaller dense loop, which then rises and interacts with the overlying field (image courtesy of Lyndsay Fletcher).]]<br />
<br />
We estimate that a column density of up to 2×10<sup>20</sup> cm-2 could have been evaporated during the initial flare. If only a fraction of this plasma is transferred onto the small dense loop, then the this may account for the already-high column density at the onset of the second flare.<br />
<br />
==Conclusions==<br />
<br />
We propose that plasma evaporated in an earlier flare may account for the already high plasma density at the onset of a second flare. The usual two-dimensional flare model runs into some trouble here as within this framework the electrons accelerated in the later flare would have almost no access to the dense loop created by an earlier event. However using multiwavelength observations to build up a three-dimensional picture of the magnetic configuration in which the flares occur, we see that reconnection may produce a short sheared loop onto which plasma evaporates. This loop then becomes unstable, allowing a second flare in which the density increases further and few accelerated electrons reach the footpoints.<br />
<br />
'''Biographical note''': Laura Bone is a PhD student at the University of Glasgow. This work was carried out in collaboration with John Brown, Lyndsay Fletcher, Astrid Veronig and Stephen White.<br />
<br />
[[Has observation by:: RHESSI| ]]</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Birth_of_a_dense_flaring_loopBirth of a dense flaring loop2018-09-17T18:59:55Z<p>Schriste: /* SOL2002-04-14 */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Birth of a dense flaring loop<br />
|number=23<br />
|first_author=Laura Bone<br />
|publish_date=27 January 2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::24]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::22]]}}<br />
}}<br />
<br />
<br />
==Introduction==<br />
<br />
In the classical "thick target" model of a solar flare, high-energy electrons are accelerated at some reconnection site in the corona. The electrons stream down the magnetic field lines to the chromosphere where they are stopped by the dense plasma and produce hard X-ray emission by collisional bremsstrahlung. The plasma at the footpoints then heats up, and "evaporates" into the coronal magnetic loop, producing thermal soft X-ray emission.<br />
<br />
From the early 1970's hard X-ray emission has occasionally been seen in the corona as well, initially in occulted events (flares with the footpoints behind the limb). These first events were interpreted to be thermal bremsstrahlung from hot coronal plasma, with significant thick target emission only at the footpoints. The first observation of an impulsive, and by implication non-thermal coronal source, was the famous <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/An_Alternative_View_of_the_Masuda_Flare Masuda flare]. <br />
Later Yohkoh studies of X-ray flares demonstrated that coronal sources exhibit both gradual and impulsive hard X-ray emission.<br />
<br />
Since its launch in 2002 RHESSI, has observed a number of both moving and stationary coronal sources. The frequency of these sources and the diversity of the physics they show us has been one of the major developments of the RHESSI era. Theoretically one of the simplest explanations for coronal sources, particularly in flares with little or no footpoint emission, is an enhanced coronal density, but then we must explain how the density came to be so large in the first place. We must also ask whether these dense regions can be observed in other ways, e.g. via the techniques of solar radio astronomy.<br />
<br />
== SOL2002-04-14 ==<br />
<br />
We have studied a flare in which RHESSI observes emission up to 50 keV, definitely in the "hard" or "non-thermal" energy range. This unocculted M3.7 flare, on the [[Has event date:: April 14, 2002 00:05]]<br />
, showed hard X-ray emission primarily from the corona, with little or no footpoint emission during most of the event (see Figure 1). This was interpreted in terms of high coronal column densities: electrons up to ~50 keV actually stopped high up in the corona, in effect producing a coronal thick target.<br />
<br />
[[File:23f1.png|600px|thumb|center|Figure 1: RHESSI images at 6-12keV with overlayed contours at 25-50 keV. Footpoint emission is only visible in the higher energy band for a brief period during the flare maximum, around 00:10 UT. During the rest of the flare the X-ray emission is confined primarily to the loop top (image courtesy of Astrid Veronig).]]<br />
<br />
At the onset of this event emission measures derived from the GOES soft X-ray data suggested that the column density was already rather high, ~1020cm-2, implying that the loop was already filled with plasma. For this flare complementary multi-wavelength observations provided us with information allowing us to understand the active region evolution and sequence of events in which the dense loop was born. Two flares occurred in this active region in close temporal and spatial proximity and this sequence of flares leads us to propose a series of reconnections between multiple loops in a 3-dimensional field geometry.<br />
<br />
==Model of the loop interaction==<br />
<br />
The loop in which the M3.7 flare occurs can be seen to brighten gradually both in EIT 195 Å and 17 GHz (radio microwave band) just after the impulsive phase of the earlier flare. Combined with the already bright, polarised radio emission from the southern footpoint, this is evidence that fast electrons may already have been present in the flaring loop (see Figure 2).<br />
<br />
[[File:23f2.png|600px|thumb|center|Figure 2: The evolution of the microwave emission (grayscale images) over the course of the flare; contours show radio polarization. The first panel shows the latter half of the earlier flare. The second image is from between the flares. The third panel shows the whole loop in which the M3.7 flare occurs beginning to brighten and the fourth image shows the whole loop at the peak of the flare (image courtesy of Stephen White).]]<br />
<br />
The earlier flare occurs at the interface of two active regions through reconnection it forms a small dense loop between the two active regions and an area of overlying field. The small loop, observed to be initially rather dense, becomes unstable; this leads to further particle acceleration and density increase and the rise of the loop which interacts with the overlying field (see Figure 3).<br />
<br />
[[File:23f3.png|600px|thumb|center|Figure 3: The sequences of reconnection occurring between the magnetic loops in the two flares. The initial reconnection forms the overlying loops and the smaller dense loop, which then rises and interacts with the overlying field (image courtesy of Lyndsay Fletcher).]]<br />
<br />
We estimate that a column density of up to 2×10<sup>20</sup> cm-2 could have been evaporated during the initial flare. If only a fraction of this plasma is transferred onto the small dense loop, then the this may account for the already-high column density at the onset of the second flare.<br />
<br />
==Conclusions==<br />
<br />
We propose that plasma evaporated in an earlier flare may account for the already high plasma density at the onset of a second flare. The usual two-dimensional flare model runs into some trouble here as within this framework the electrons accelerated in the later flare would have almost no access to the dense loop created by an earlier event. However using multiwavelength observations to build up a three-dimensional picture of the magnetic configuration in which the flares occur, we see that reconnection may produce a short sheared loop onto which plasma evaporates. This loop then becomes unstable, allowing a second flare in which the density increases further and few accelerated electrons reach the footpoints.<br />
<br />
'''Biographical note''': Laura Bone is a PhD student at the University of Glasgow. This work was carried out in collaboration with John Brown, Lyndsay Fletcher, Astrid Veronig and Stephen White.</div>Schristehttps://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/Birth_of_a_dense_flaring_loopBirth of a dense flaring loop2018-09-17T18:58:32Z<p>Schriste: /* Biographical note */</p>
<hr />
<div>{{Infobox Nugget<br />
|name = Nugget<br />
|title=Birth of a dense flaring loop<br />
|number=23<br />
|first_author=Laura Bone<br />
|publish_date=27 January 2006<br />
|next_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::24]]}}<br />
|previous_nugget={{#ask: [[Category:Nugget]] [[RHESSI Nugget Index::22]]}}<br />
}}<br />
<br />
<br />
==Introduction==<br />
<br />
In the classical "thick target" model of a solar flare, high-energy electrons are accelerated at some reconnection site in the corona. The electrons stream down the magnetic field lines to the chromosphere where they are stopped by the dense plasma and produce hard X-ray emission by collisional bremsstrahlung. The plasma at the footpoints then heats up, and "evaporates" into the coronal magnetic loop, producing thermal soft X-ray emission.<br />
<br />
From the early 1970's hard X-ray emission has occasionally been seen in the corona as well, initially in occulted events (flares with the footpoints behind the limb). These first events were interpreted to be thermal bremsstrahlung from hot coronal plasma, with significant thick target emission only at the footpoints. The first observation of an impulsive, and by implication non-thermal coronal source, was the famous <br />
[http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php/An_Alternative_View_of_the_Masuda_Flare Masuda flare]. <br />
Later Yohkoh studies of X-ray flares demonstrated that coronal sources exhibit both gradual and impulsive hard X-ray emission.<br />
<br />
Since its launch in 2002 RHESSI, has observed a number of both moving and stationary coronal sources. The frequency of these sources and the diversity of the physics they show us has been one of the major developments of the RHESSI era. Theoretically one of the simplest explanations for coronal sources, particularly in flares with little or no footpoint emission, is an enhanced coronal density, but then we must explain how the density came to be so large in the first place. We must also ask whether these dense regions can be observed in other ways, e.g. via the techniques of solar radio astronomy.<br />
<br />
== SOL2002-04-14 ==<br />
<br />
We have studied a flare in which RHESSI observes emission up to 50 keV, definitely in the "hard" or "non-thermal" energy range. This unocculted M3.7 flare, on the 14th of April 2002, showed hard X-ray emission primarily from the corona, with little or no footpoint emission during most of the event (see Figure 1). This was interpreted in terms of high coronal column densities: electrons up to ~50 keV actually stopped high up in the corona, in effect producing a coronal thick target.<br />
<br />
[[File:23f1.png|600px|thumb|center|Figure 1: Rhessi images at 6-12keV with overlayed contours at 25-50 keV. Footpoint emission is only visible in the higher energy band for a brief period during the flare maximum, around 00:10 UT. During the rest of the flare the X-ray emission is confined primarily to the loop top (image courtesy of Astrid Veronig).]]<br />
<br />
At the onset of this event emission measures derived from the GOES soft X-ray data suggested that the column density was already rather high, ~1020cm-2, implying that the loop was already filled with plasma. For this flare complementary multi-wavelength observations provided us with information allowing us to understand the active region evolution and sequence of events in which the dense loop was born. Two flares occurred in this active region in close temporal and spatial proximity and this sequence of flares leads us to propose a series of reconnections between multiple loops in a 3-dimensional field geometry.<br />
<br />
==Model of the loop interaction==<br />
<br />
The loop in which the M3.7 flare occurs can be seen to brighten gradually both in EIT 195 Å and 17 GHz (radio microwave band) just after the impulsive phase of the earlier flare. Combined with the already bright, polarised radio emission from the southern footpoint, this is evidence that fast electrons may already have been present in the flaring loop (see Figure 2).<br />
<br />
[[File:23f2.png|600px|thumb|center|Figure 2: The evolution of the microwave emission (grayscale images) over the course of the flare; contours show radio polarization. The first panel shows the latter half of the earlier flare. The second image is from between the flares. The third panel shows the whole loop in which the M3.7 flare occurs beginning to brighten and the fourth image shows the whole loop at the peak of the flare (image courtesy of Stephen White).]]<br />
<br />
The earlier flare occurs at the interface of two active regions through reconnection it forms a small dense loop between the two active regions and an area of overlying field. The small loop, observed to be initially rather dense, becomes unstable; this leads to further particle acceleration and density increase and the rise of the loop which interacts with the overlying field (see Figure 3).<br />
<br />
[[File:23f3.png|600px|thumb|center|Figure 3: The sequences of reconnection occurring between the magnetic loops in the two flares. The initial reconnection forms the overlying loops and the smaller dense loop, which then rises and interacts with the overlying field (image courtesy of Lyndsay Fletcher).]]<br />
<br />
We estimate that a column density of up to 2×10<sup>20</sup> cm-2 could have been evaporated during the initial flare. If only a fraction of this plasma is transferred onto the small dense loop, then the this may account for the already-high column density at the onset of the second flare.<br />
<br />
==Conclusions==<br />
<br />
We propose that plasma evaporated in an earlier flare may account for the already high plasma density at the onset of a second flare. The usual two-dimensional flare model runs into some trouble here as within this framework the electrons accelerated in the later flare would have almost no access to the dense loop created by an earlier event. However using multiwavelength observations to build up a three-dimensional picture of the magnetic configuration in which the flares occur, we see that reconnection may produce a short sheared loop onto which plasma evaporates. This loop then becomes unstable, allowing a second flare in which the density increases further and few accelerated electrons reach the footpoints.<br />
<br />
'''Biographical note''': Laura Bone is a PhD student at the University of Glasgow. This work was carried out in collaboration with John Brown, Lyndsay Fletcher, Astrid Veronig and Stephen White.</div>Schriste