Atmosphere and Magnetosphere of the Sun

THE CHROMOSPHERE

Above the photosphere lies the cooler chromosphere, the inner part of the solar atmosphere. This region emits very little light of its own and cannot be observed visually under normal conditions. The photosphere is just too bright, dominating the chromosphere's radiation. The relative dimness of the chromosphere results from its low density--large numbers of photons simply cannot be emitted by a tenuous gas containing very few atoms per unit volume. Still, although it is not normally seen, astronomers have long been aware of the chromosphere's existence. Figure 16.10 shows the Sun during an eclipse in which the photosphere--but not the chromosphere--is obscured by the Moon. The chromosphere's characteristic reddish hue is plainly visible. This coloration is due to the H ("hydrogen alpha") emission line of hydrogen, which dominates the chromospheric spectrum. (Recall from Chapter 4 that the wavelength of this line is 656.3 nm, or 6563 , right in the middle of the red portion of the spectrum.)

Figure 16.10 This photograph of a total solar eclipse shows the solar chromosphere, a few thousand kilometers above the Sun's surface.

The chromosphere is far from tranquil. Every few minutes, small solar storms erupt, expelling jets of hot matter known as spicules into the Sun's upper atmosphere (see Figure 16.11). These long, thin spikes of matter leave the Sun's surface at typical velocities of about 100 km/s, reaching several thousand kilometers above the photosphere. Spicules are not spread evenly across the solar surface. Instead, they cover only about 1 percent of the total area, tending to accumulate around the edges of supergranules. The Sun's magnetic field is also known to be somewhat stronger than average in those regions. Scientists speculate that the downwelling material there tends to strengthen the solar magnetic field and that spicules are the result of magnetic disturbances in the Sun's churning outer layers.

Figure 16.11 Solar spicules, short-lived narrow jets of gas that typically last mere minutes, can be seen sprouting up from the solar chromosphere in this H image of the Sun. The spicules are the thin, dark, spikelike regions. They appear dark against the face of the Sun because they are cooler than the solar photosphere.

THE CORONA

During the brief moments of an eclipse, if the Moon's angular size is large enough that both the photosphere and the chromosphere are blocked, the ghostly solar corona can be seen, as in Figure 16.12. With the photospheric light removed, the pattern of spectral lines changes dramatically. The intensities of the usual lines alter, suggesting changes in elemental abundances or gas temperature or both. The spectrum shifts from absorption to emission, and an entirely new set of spectral lines suddenly appears. These new coronal (and in some cases chromospheric) lines were first observed during eclipses in the 1920s. For years afterward, some researchers (for want of any better explanation) attributed them to a nonterrestrial element, which they called "coronium."

Figure 16.12 When both the photosphere and the chromosphere are obscured by the Moon during a solar eclipse, the faint corona becomes visible. This photograph shows clearly the emission of radiation from the solar corona.

We now recognize that these new spectral lines do not indicate any new kind of atom. Coronium does not exist. Rather, the new lines arise because atoms in the corona have lost several more electrons than atoms in the photosphere--that is, the coronal atoms are much more highly ionized. For example, astronomers have identified coronal lines corresponding to iron ions with as many as 13 of their normal 26 electrons missing. In the photosphere, most iron atoms have lost only 1 or 2 of their electrons. The cause of this extensive electron stripping is the high coronal temperature. The degree of ionization inferred from spectra observed during solar eclipses tell us that the gas temperature of the upper chromosphere exceeds that of the photosphere. Furthermore, the temperature of the solar corona, where even more ionization is seen, is still higher.

Based on many observations of conditions at different distances from the limb of the Sun, from the photosphere outward into the corona, Figure 16.13 plots the variation of gas temperature with altitude. The temperature decreases to a minimum of about 4500 K some 500 km above the photosphere, after which it rises steadily. About 1500 km above the photosphere, the gas temperature begins to rise rapidly, reaching more than 1,000,000 K at an altitude of 10,000 km. Thereafter, it remains roughly constant. Using this temperature profile, we can draw a clear distinction between the chromosphere and the corona: The chromosphere extends from the top of the photosphere for approximately 1500 km. The region in which the temperature rises rapidly--from about 1500 km to 10,000 km--is called the transition zone. At 10,000 km, the corona begins.

Figure 16.13 The change of gas temperature in the lower solar atmosphere is dramatic. The minimum temperature marks the outer edge of the chromosphere. Beyond that, the temperature rises sharply in the transition zone, finally leveling off at over 1,000,000 K in the corona.

The cause of this rapid temperature rise is not fully understood. The temperature profile runs contrary to intuition--moving away from a heat source, we would normally expect the heat to diminish, but this is not the case for the lower atmosphere of the Sun. The corona must have another energy source. Astronomers now believe that magnetic disturbances in the solar photosphere--a little like spicules, but on a much larger scale--are ultimately responsible for heating the corona. We will return to these disturbances in more detail in the next section.

THE SOLAR WIND

Electromagnetic radiation and fast-moving particles--mostly protons and electrons--escape from the Sun all the time. The radiation moves away from the photosphere at the speed of light, taking 8 minutes to reach Earth. The particles travel more slowly, although at the still considerable speed of about 500 km/s, reaching the Earth in a few days. This constant stream of escaping solar particles is the solar wind.

The solar wind results from the high temperature of the solar corona. About 10,000,000 km above the photosphere, the coronal gas is hot enough to escape the Sun's gravity, and it begins to flow outward into space. At the same time, the solar atmosphere is continually replenished from below. If that were not the case, the corona would disappear in about a day. The Sun is, in effect, "evaporating"--constantly shedding mass through the solar wind. But the wind is an extremely thin medium. Although it carries away about a million tons of solar matter each second, less than 0.1 percent of the Sun has been lost since the solar system formed billions of years ago. Our star is indeed evaporating, but it is losing only a negligible fraction of its huge bulk.

THE SUN IN X RAYS

What sort of radiation is emitted by a gas of 1,000,000 K? Unlike the 6000 K photosphere, which emits most strongly in the visible part of the electromagnetic spectrum, the hotter coronal gas radiates at much higher frequencies--primarily in X rays. For this reason, X-ray telescopes have become important tools in the study of the solar corona. Figure 16.14 shows an X-ray image of the Sun. The full corona extends well beyond the regions shown, but the density of coronal particles emitting the radiation diminishes rapidly with distance from the Sun. The intensity of X-ray radiation farther out is too dim to be seen here.

 

Figure 16.14 Images of X-ray emission from the Sun observed by the Skylab space station. These frames were taken at one-day intervals. Note the dark, boot-shaped coronal hole traveling from left to right, where the X-ray observations outline in dramatic detail the abnormally thin regions through which the high-speed solar wind streams forth.

In the mid-1970s, instruments aboard NASA's Skylab space station revealed that the solar wind escapes mostly through solar "windows" called coronal holes. The dark area moving from left to right in Figure 16.14 represents a coronal hole. Not really holes, such structures are simply deficient in matter--vast regions of the Sun's atmosphere where the density is about 10 times lower than the already tenuous, normal corona. Coronal holes are underabundant in matter because the gas there is able to stream freely into space at particularly high speeds, driven by disturbances in the Sun's atmosphere and magnetic field. In coronal holes, the solar magnetic field lines extend from the surface far out into interplanetary space. Charged particles tend to follow the field lines, so they can escape. In other regions of the corona, the solar magnetic field lines stay close to the Sun, keeping charged particles near the surface and inhibiting the outward flow of the solar wind (just as Earth's magnetic field tends to prevent the incoming solar wind from striking Earth), and the density remains (relatively) high. The largest coronal holes can be hundreds of thousands of kilometers across. Structures of this size are seen only a few times each decade. Smaller holes--perhaps only a few tens of thousand kilometers in size--are much more common, appearing every few hours.

 

 Supplement

In 1962, the Mariner 2 probe confirmed the suspicions of the astronomers of the XIXth century: a "solar wind" permanently blows from the Sun to the confines of the solar system and delimits its "area of influence", the heliosphere. It consists of a plasma (ionized gas) which mainly contains protons and electrons. Although corresponding to a loss of mass of a million tons a second, this plasma is far from dense. It nevertheless conveys the magnetic field of the Sun, which thus arrives to the Earth. ULYSSES probe revealed that in fact there are two kinds of wind: a "fast" wind (700 km/s) which especially escapes from the poles through coronal holes, areas where the magnetic field of the Sun is less intense, and a "slow" wind (400 km/s), mainly in the low heliographic latitudes, where the constraint of the magnetic field is stronger.

Our star has another characteristic: has it is not a rigid body, its different regions do not rotate at the same speed at the poles and the equator. This differential rotation makes the lines of magnetic field twist, break and reconnect, forming loops. Sometimes, these ones, filled with plasma, break in a brutal release of energy in the form of radiation (, X, UV) and of particles. The hot gas is expelled in the heliosphere. Then it is a flare. If it occurs on the side of the Sun facing the Earth, the flood of very fast particles (faster than 30.000 km/s) hardly takes 30 to 60 minutes to travel the 150 million kilometers which separate us from the Sun.

Higher, in the corona - a very wide area constituting the external atmosphere of the Sun (it is the very thin silver plated veil which can be seen durint solar eclipses) -, similar phenomena can occur: coronal matter ejections, commonly called CME (Coronal Mass Ejections). These disturbances of solar wind (probably due to magnetic phenomena) producing an expulsion of plasma and a shock wave which propagates in the heliosphere. Ejected particles, although slower than the ones from the flares, arrive to the Earth within two or three days.