Coronal implosion

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Number: 95
1st Author: Rui Liu
2nd Author:
Published: 16 February 2009
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During solar flares and coronal mass ejections (CMEs), magnetic reconnection results in the rapid release of free energy stored in the coronal magnetic field. Specifically, the magnetic free energy is believed to be stored in the "nonpotential" (sheared or twisted) magnetic field of a filament channel. Eruptions occur if the force balance between the upward magnetic pressure force of the sheared filament channel field, and the downward magnetic tension force of the overlying quasi-potential field, is disrupted. As Hudson (2000) pointed out, a reduction of magnetic energy leads to a reduction of the upward magnetic pressure, which would inevitably result in the contraction of the overlying field. However, the large-scale motions observed in almost all flares or CMEs are always explosive rather than implosive, which poses an obvious observational dilemma.

On the other hand, observations of contracting flaring loops during the early phase of flares, manifested by converging conjugate footpoints and descending looptop emission, have been reported in X-ray, EUV, H-alpha, and microwave observations. The contraction can be explained by the relaxation of the sheared magnetic field, a scenario consistent with the Hudson conjecture.

In this Nugget, observations of a coronal implosion are presented, showing that the EUV coronal loops overlying an eruptive filament push inward during a flare, which is associated with the converging motion of the conjugate HXR footpoints and the downward motion of the HXR looptop source.


The (6.7MB QuickTime) movie shows a collection of criss-crossed coronal loops (labeled 'L' in Figure 1) that experienced contraction during a GOES class C8.9 flare on 2005 July 30. These loops were overlying a dark filament (labeled 'F' in Figure 1). The flare exhibits a relatively long pre-heating phase (Figure 2 (third panel)), during which RHESSI hard X-ray (HXR) count rates at lower energies (below 25 keV) as well as GOES soft X-ray (SXR) fluxes began to increase gradually as early as 16:40:00 UT, and the flare emission is dominated by a thermal looptop source. The impulsive phase started over 6 min later from 16:46:36 UT onward, during which impulsive HXR bursts (above 25 keV) are observed, and the flare emission is composed of a pair of conjugate footpoints in addition to the looptop emission.

Figure 1 TRACE 171 Å image prior to a GOES C8.9 flare, overlaid with centroid positions of the HXR looptop emission (in a '+' symbol) and of the HXR footpoints (in a 'x' symbol), observed by RHESSI during the flare. The time evolution of the centroid positions is color-coded, and the positions are differentially rotated to 2005 July 30 16:40:34 UT, the time when the TRACE image was taken. A horizontal slit is placed across the group of coronal loops of interest (labeled 'L') to study their contraction. Underneath, a filament appearing in absorption is labeled 'F'.

From Figure 1 and Figure 2 one can see that the collection of coronal loops of interest is mainly composed of three clusters of loops at different altitudes. In Figure 2 (top panel) all three clusters of loops contracted at approximately the same speed (4 - 7 km s-1), starting at approximately the same time (~ 16:42 UT). The higher and middle clusters of loops appear to expand first (~ 16:48 UT) while the lower cluster of loops were still contracting till about 16:53 UT, well into the gradual phase of the flare, and then began to expand at about 60 km s-1. The whole structure, including the filament underneath, was observed to erupt in TRACE 171 Å at about 17:02 UT, and resulted in a fast CME. The contraction of the overlying coronal loops is associated with the converging motion of the conjugate HXR footpoints during the impulsive phase, as well as the downward motion of the HXR looptop source toward the solar surface as early as the pre-heating phase (Figure 1 and Figure 2 (middle panel)).

Figure 2 Coronal implosion observed in EUV and X-rays. a) CME height observed in LASCO C2 and C3. Hatched area indicates the time interval covered by (b-d). b) Slices of TRACE 171 Å images cut by the horizontal slit as defined in Figure 1 are rotated to a vertical direction and placed on the time axis. The slice width on the time axis reflects the exposure duration of each image and voids are due to data gaps. c) Time evolution of the height of the HXR looptop emission (in a '+' symbol) as well as the separation between two conjugate HXR footpoints (in an 'x' symbol). The "height" of the looptop HXR source is defined as the perpendicular distance of its centroid position to the line connecting the centroid positions of the conjugate footpoints at 16:47:08 UT. c) RHESSI count rates (s-1) in the 6 - 15 keV (grey) and 25 - 50 keV (black) energy ranges (scaled by the y-axis on the left), and GOES 1 - 8 Å (0.5 - 4 Å) fluxes in dotted (dashed) line (scaled by the y-axis on the right). Note that the sudden transition in the 6 - 15 keV lightcurve was due to the change of the RHESSI attenuator state from A0 to A1 at 16:44:08 UT, and that the observation was discontinued at 16:49:28 UT when RHESSI went into eclipse.

Why Is Coronal Implosion So Rare?

When a flare has not induced any significant change to the confining field, a decrease of the magnetic pressure in the flaring region could be compensated by an increase of the thermal pressure in the flaring region (a conservative estimation of the plasma β at the flare looptop in this observation is as large as 0.7). This makes sense due to the possible localized heating of the coronal plasma by the energy release and also due to chromospheric evaporation which injects hot plasma into the corona. On the other hand, if the overlying field lines which provide the confinement undergo breakout reconnection, or, are stretched and thereby reconnected beneath a rising and expanding fluxrope as in the standard flare model, the reduction of the magnetic tension force may be comparable to or greater than the reduction of the magnetic pressure force. Hence no implosion.

Conclusions (What Is the Role of Pre-heating?)

During the pre-heating phase (which may correspond to a slow-reconnection stage), the flare loop density increases due to "gentle" chromospheric evaporation in response to the hot, dominant coronal source (20 - 30 MK). At the onset of the impulsive phase (corresponding to a fast-reconnection stage), the coronal atmosphere is already heated to a dense and hot state. Hence the column depth of the corona is enhanced and the stopping distance for electrons is shorter. Therefore, the bulk of the accelerated electrons will lose their energy in the corona, and chromospheric evaporation will be strongly suppressed, due not only to reduced energy deposited in the chromosphere, but also to the increased inertia of the overlying material.

As a result, the increase in thermal pressure is probably not enough to counterbalance the decrease of magnetic pressure, because the released free energy can escape in other forms, such as optically-thin radiation, thermal conduction, and hydromagnetic waves. Thus, the magnetic structure surrounding the flaring region, observed as the EUV coronal loops, pushes inward. For the overlying loops to “know” the reduction of the magnetic pressure in the reconnection region, however, the pre-heating phase should last for at least t = L / VA ~ 50 s where L is the height of the overlying loops with respect to the reconnection location (~5x109 cm) and VA is the Alfvén speed in the corona (~1000 km/s).

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