About the Sun I

About the Sun I

The Sun is the nearest star and provides the energy necessary to sustain life here on earth. It also provides a means to study the physics of magnetic plasmas in an environment that cannot be reproduced in a terrestrial laboratory. Our Sun is a main sequence, variable star --- essentially a large (1.4 million kilometers in diameter) massive ball of gas that is dense enough at its core to allow thermonuclear reactions to occur. The outward pressure from this continuous nuclear explosion is held in check by the mutual gravitational interaction of atoms that act to collapse the star inward, thus providing a stable pressure equilibrium for the entire system.

The nuclear fusion at the Sun's core produces an enormous amount of energy that must be transported through the rest of the solar interior before it escapes through the solar atmosphere as radiation. The transport of energy occurs in different ways in different regions of the interior (see Figure 1): above the energy generating core, energy transport occurs radiatively; photons are absorbed and re-emitted (in a random direction) by atoms of the dense hot plasma. Thus, they "random-walk" their way through the interior (a process that can take hundreds of thousands of years). As photons reach the outer 30% of the solar envelope (where it is somewhat cooler, yet still dense), the atoms aren't so willing to immediately re-emit the photons into the surroundings, thus rendering the plasma opaque, and slowing the transport of energy via radiation.

In this region, the most efficient mechanism of energy transport is convection (hence the upper 30% of the solar interior is dubbed "the convection zone"). Here, hot plasma (relative to its surroundings) will buoyantly rise toward the less-dense surface layers, where it cools and sinks back down toward the base of the convection zone, creating a turbulent, rolling motion reminiscent of that observed in a pot of boiling water. The net transport of energy is outward, and occurs much more quickly than in the radiative interior. Figure 2 shows a signature of convection: the characteristic granulation pattern at the visible surface of the Sun.

At the top of the convection zone lies the visible surface called the photosphere. It is within this thin rippled layer that the plasma becomes transparent to optical continuum radiation ("white light"), allowing a majority of the electromagnetic energy reaching this layer to escape into space (to be seen only ~8.3 minutes later at earth). While continuum radiation escapes, the photosphere remains optically thick to spectral lines --- radiation that results from changes in the energy levels of electrons bound to atoms. Above the photosphere lies the chromosphere; a complex layer transparent to the optical continuum, yet still optically thick to strong spectral lines. Higher yet, lies the solar corona (see Figure 3) --- the hot optically thin plasma that is visible during a total solar eclipse. The combined layers of the photosphere, chromosphere, corona (and the complex transition between the cool, roughly fourty five hundred degree K chromosphere and the hot, 1-2 million degree K corona) is referred to as the solar atmosphere.

One active area of research in solar physics is to determine exactly why the corona is so hot --- after all, one would expect that the farther away from the core the plasma is, the cooler it should become (heat flowing from a cool region to a hot region appears to violate the second law of thermodynamics!). Thus, some sort of a heating mechanism must exist in order for the tenuous corona to maintain such a high temperature. The exact nature of coronal heating is still a topic of debate, but most scientists agree that the Sun's magnetic field plays an important role in this process.