Solar atmosphere

Historically, the visible emission lines shown in Figure 15-3 were the first atomic hydrogen lines discovered. They were found in the spectrum of the sun by W. H. Wollaston in 1802. In 1862, A. J. Angstrom announced that there must be hydrogen in the solar atmosphere. These lines were detected first because of the lesser experimental difficulties in the visible spectral region. They are called the "Balmer series because J. J. Balmer was able to formulate a simple mathematical relation among the frequencies (in It S). The ultraviolet series shown in Figure 15-3 was... 

Fig. 25-3. Some molecules detected in the solar atmosphere and the elements they contain.

Room-temperature fluorescence (RTF) has been used to determine the emission characteristics of a wide variety of materials relative to the wavelengths of several Fraunhofer lines. Fraunhofer lines are bands of reduced intensity in the solar spectrum caused by the selective absorption of light by gaseous elements in the solar atmosphere. RTF studies have recently included the search for the causes of the luminescence of materials and a compilation of information that will lead to "luminescence signatures" for these materials. For this purpose, excitation-emission matrix (EEM) data are now being collected. 

Fig. 3.7. Dependence of the line absorption coefficient on distance from the line centre, for the particular case of the Na I D-lines in the outer solar atmosphere. After Unsold (1977). Copyright by Springer-Verlag New York Inc.

Note that iron is sufficiently ionized in the solar atmosphere that the abundance of Fe I can be neglected and its partition function (or the ground-state statistical weight gi) and the electron pressure cancel out. 

Bernard Ephraim Julius Pagel was bom in Berlin on 4 January 1930, but when his father was dismissed from his post as Jewish persecution increased, the family moved to Britain in 1933. From Merchant Taylors School he won an open scholarship in Natural Sciences at Sidney Sussex College, Cambridge, graduating with First-class honours in Physics in 1950. His early research at Cambridge (Ph.D. 1955) centred on the solar atmosphere. Inspired by Willy Fowler, a future Nobel Prize winner who was visiting from California, he started a life-long interest in the abundances of the chemical elements. 

The solar atmosphere is too cool to lead to helium excitation, and hence helium lines, except in exceptional conditions associated with flares. The helium content of the Sun is therefore inferred from observations of hotter stars. 

The cross section obtained for single fullerenes and buckyonions reproduce the behaviour of the interstellar medium UV extinction curve. A power-law size distribution n(R) R m with in = 3.5 1.0 for these molecules can explain the position and widths observed for the 2,175 A bump and, partly, the rise in the extinction curve at higher energies. We infer ISM densities of 0.2 and 0.1 ppm for small fullerenes and buckyonions (very similar to the densities measured in meteorites). If as expected the cosmic carbon abundance is close to the solar atmosphere value, individual fullerenes may lock up 20-25% of the total carbon in the diffuse interstellar space. 

B. Stromgren, On the chemical composition of the solar atmosphere, 218-257 in Festschrift fir Elis Stromgren (Copenhagen, 1940), on 257. See also B. Stromgren, Scientists I have known and some astronomical problems I have met," Annual Review of Astronomy and Astrophysics 21 (1983) 1-11. 

This figure shows the match between the bright lines in the spectrum of iron with some of the Fraunhofer lines, proving the presence of iron in the solar atmosphere. 

About 40 percent of the solar granules explode, vomiting their material into the solar atmosphere at speeds of 1,000 miles per hour (or about 500 meters per second). If we could hear sound waves across the vacuum of space that separates us from the Sun, we would hear a thunderous roar accompanying each granulation eruption. 

Depending on the rate at which the Sun loses mass on its way to becoming a white dwarf, there are different fates for the Earth. If mass is lost early on, the dry Earth could escape being vaporized as its orbit expands in response to the Sun s lower mass. As a result, the Earth would orbit near Mars s current position as a cenotaph to humankind s past glory. It might be a featureless cenotaph, and I wonder if we could trace some of Earth s ancient shorelines and ocean basins. On the other hand, if the mass loss occurs late in the Sun s evolution, the Earth may actually orbit inside the outer solar atmospheres. In this case, the Earth would be burnt to a crisp and pulled deeper into the Sun. 

From Eq. (4.3), wo see that as the pressure is reduced at constant temperature, the dissociation becomes greater, until finally at vanishing pressure the dissociation can become complete, even at ordinary temperatures. This is a result of importance in astrophysics, as has been pointed out by Saha. In the solar atmosphere, there is spectroscopic evidence of the existence of rather highly ionized elements, even though the temperature of the outer layers of the atmosphere is not high enough for us to expect such ionization, at ordinary pressures. However, the pressure in these layers of the sun is extremely small, and for that reason the ionization is abnormally high. 

Calculation of the cross-sections for the photodissociation of MgH in the solar atmosphere... 

Models of planetary evolution assume that at the time of planetary formation the solar system had a single universal and well-mixed composition from which aU parts of the solar system were derived (see Podosek, 1978). Information as to the elemental and isotopic characteristics of this primordial composition is presently available from the Sun, meteorites, and the atmospheres of the giant planets (Wider, 2002). In the case of the Sun, distinction is usually made between the present-day composition, which is available via spectral analysis of the solar atmosphere and capture of the solar wind, either directly in space or by using metallic foU targets, and the proto-Sun (the composition at the time of planetary accretion) whereby the lunar regolith and/or meteorites are utilized as archives of ancient solar wind. As discussed below, the distinction is only really important for helium due to production of He by deuterium burning. 

The relative abundance of the more common elements in the solar atmosphere is shown in Figure 2 and Table 2. The remainder of the chemical elements occurs only in trace quantities. By any chemist s measure, the universe is composed mainly of hydrogen and helium, with some impurities of practically all elements. 

Table 2. Relative abundance of the more common elements in the solar atmosphere (after Cameron, 1982).

This model might also explain the observed depletion of neon in the planet s atmosphere. The Galileo probe s Neutral Mass Spectrometer (GPMS) observed an abundance of neon about one tenth that in the solar atmosphere. Perhaps, researchers hypothesize, neon dissolves in helium droplets as they form in the atmosphere and fall into the planet s interior. This phenomenon is possible because helium and neon, unlike helium and hydrogen, are completely miscible. 

Converting the absorption lines into abundances requires knowledge of line positions of neutral and ionized atoms, as well as their transition probabilities and lifetimes of the excited atomic states. In addition, a model of the solar atmosphere is needed. In the past years, atomic properties have seen many experimental updates, especially for the rare earth elements (see below). Older solar atmospheric models used local thermodynamic equilibrium (LTE) to describe the population of the quantum states of neutral and ionized atoms and molecules according to the Boltzmann and Saha equations. However, the ionization and excitation temperatures describing the state of the gas in a photospheric layer may not be identical as required for LTE. Models that include the deviations from LTE (=non-LTE) are used more frequently, and deviations from LTE are modeled by including treatments for radiative and collision processes (see, e.g., [27,28]). 

Solar atmospheric models have evolved from one dimension to more complicated 2D and 3D models designed to take into account effects of convection and granulation on radiative transfer in the solar atmosphere. Recent applications of 3D models instead of older ID models leads to lower abundances of several elements, notably oxygen, from previously determined values. Significant reductions of photospheric abundances in other elements (e.g., Na, Al, Si) were also found (see, e.g., [15]). However, different 3D model assumptions lead to different results, see for example the discussion by [29] for silicon. 

Lower abundances for some elements derived from these models also cause problems for standard solar models that describe the evolution of the sun to its current radius and luminosity (see [30]). Another problem is that some of the 3D abundances compare worse to meteoritic data than before. In the following preference is given to elemental abundances derived with more conservative solar atmospheric models however, for some elements (e.g., P, S, Eu) the results from 1D/2D and some 3D models produce consistent results. 

The production of radiation by the Sun depends strongly on the physical and chemical structure of the solar atmosphere. Therefore, in order to understand the nature and variations of the radiation incident on the Earth, we briefly present some of the important aspects of the Sun s atmosphere, whose structure is schematically represented in Figures 4.5 and 4.6. 

Figure 4-5. Schematic diagram of the solar atmosphere with its different layers and physical characteristics.

---