Astronomers can glean an enormous amount of information about the Sun from an analysis of the absorption lines that arise in the photosphere and lower atmosphere. 
(Sec. 4.4) Figure 16.8 (see also Figure 4.4) is a detailed spectrum of the Sun, obtained for a small portion of the range of visual wavelengths, from 360 to 690 nm. Notice the intricate dark Fraunhofer absorption lines superposed on the background continuous spectrum.
Figure 16.8 A detailed spectrum of our Sun in a portion of the visible domain shows thousands of Fraunhofer spectral lines, which indicate the presence of some 67 different elements in various stages of excitation and ionization in the lower solar atmosphere.
As discussed in Chapter 4, spectral lines arise when electrons in atoms or ions make transitions between states of well-defined energies, emitting or absorbing photons of specific energies (that is, wavelengths or colors) in the process. (Sec. 4.2) However, to explain the spectrum of the Sun (and, indeed, the spectra of all stars), we must modify slightly our earlier description of the formation of absorption lines. We explained these lines in terms of cool foreground gas intercepting light from a hot background source. In actuality, both the bright background and the dark absorption lines in Figure 16.8 form in roughly the same location in the Sun—the solar photosphere and lower chromosphere. To understand how line formation occurs, let's reconsider the solar energy-emission process in a little more detail.
Below the photosphere, the solar gas is sufficiently dense, and interactions among photons, electrons, and ions sufficiently common, that radiation cannot escape directly into space. In the solar atmosphere, however, the probability that a photon will escape without further interaction with matter depends on its energy. If that energy happens to correspond to some electronic transition in one of the atoms or ions present in the gas, then the photon may be absorbed again before it can travel very far—the more elements present of the type suitable for absorption, the lower the escape probability. Conversely, if the photon's energy does not coincide with any such transition, then the photon cannot interact further with the gas, and it leaves the Sun headed for interstellar space, or perhaps the detector of an astronomer on Earth.
Thus, as illustrated in Figure 16.9, when we look at the Sun, we are actually peering down into the solar atmosphere to a depth that depends on the wavelength of the light under consideration. Photons with wavelengths well away from any absorption feature tend to come from deep in the photosphere, whereas those at the centers of absorption lines, being much more likely to interact with matter as they travel through the solar gas, mainly escape from higher, cooler levels. The lines are darker than their surroundings because the temperature at the level of the atmosphere where they form is lower than the 5800 K temperature at the base of the photosphere, where most of the continuous emission originates. (Recall that by Stefan's law, the brightness of a radiating object depends on its temperature—the cooler the gas, the less energy it radiates.) (Sec. 3.4) The existence of the Fraunhofer lines is direct evidence that the temperature in the Sun's atmosphere decreases with height above the photosphere.
Figure 16.9 Formation of solar absorption lines. Photons with energies well away from any atomic transition can escape from relatively deep in the photosphere, but those with energies close to a transition are more likely to be reabsorbed before escaping, so the ones we see on Earth tend to come from higher, cooler levels in the solar atmosphere. The inset shows a close-up tracing of two of the thousands of solar absorption lines, those produced by calcium at about 395 nm.
Tens of thousands of spectral lines have been observed and cataloged in the solar spectrum. In all, some 67 elements have been identified in the Sun, in various states of ionization and excitation. (Sec. 4.2) More elements probably exist there, but they are present in such small quantities that our instruments are simply not sensitive enough to detect them. Table 16.2 lists the 10 most common elements in the Sun. Notice that hydrogen is by far the most abundant element, followed by helium. This distribution is just what we saw on the jovian planets, and it is what we will find for the universe as a whole.
| TABLE 16.2 The Composition of the Sun | |||||||||||||||||||||||||||||||||||
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Strictly speaking, spectral analysis allows us to draw conclusions only about the part of the Sun where the lines form—the photosphere and chromosphere. However, most astronomers believe that, with the exception of the solar core (where nuclear reactions are steadily changing the composition—see Section 16.5), the data in Table 16.2 are representative of the entire Sun.
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 red H 
(hydrogen alpha) emission line of hydrogen, which dominates the chromospheric spectrum. (Sec. 4.2)
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 (Figure 16.11). These long, thin spikes of matter leave the Sun's surface at typical speeds 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 Himage 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.
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 (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 composition or 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 new nonterrestrial element, which they dubbed "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.
Coronal Ejection and Grazing Comet
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. Therefore, their internal electronic structures, and hence their spectra, are quite different from the structure and spectra of atoms and ions in the photosphere. 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 tells 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 higher still.
Figure 16.13 shows how the temperature of the Sun's atmosphere varies 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.
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 (Section 16.4).
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 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 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 continuously 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. The wind is an extremely thin medium, however. 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 4.6 billion years ago.
What sort of radiation is emitted by a gas of 1,000,000 K? Unlike the 5800 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. (Sec. 3.4) For this reason, X-ray telescopes have become important tools in the study of the solar corona. Figure 16.14 shows X-ray images 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 1-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 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), so 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.
