Among the most extensive applications of atomic physics in astronomy is the precise computation of transfer of radiation from a source through matter. The physical problem depends in part on the bulk temperature and density of the medium through which radiation is propagating. Whether the medium is relatively transparent or opaque (‘thin’ or ‘thick’) depends not only on the temperature and the density, but also on the atomic constituents of matter interacting with the incident radiation via absorption, emission and scattering of radiation by particular atomic species in the media. Since optical lines in the visible range of the spectrum are most commonly observed, the degree of transparency or opaqueness of matter is referred to as optically thin or optically thick. However, it must be borne in mind that in general we need to ascertain radiative transfer in all wavelength ranges, not just the optical. Macroscopically, we refer to optical thickness of a whole medium, such as a stellar atmosphere. But often one may observe a particular line and attempt to ascertain whether it is optically thick or thin in traversing the entire medium.

Radiative transfer and atomic physics underpin quantitative spectroscopy. But together they assume different levels of complexity when applied to practical astrophysical situations. Significantly different treatments are adopted in models for various astrophysical media. At low densities, prevalent in the interstellar medium (ISM) or nebulae, ne <; 106 cm-3, the plasma is generally optically thin (except for some strong lines, such as the Lyα, that do saturate), and consideration of detailed radiative transfer effects is not necessary.

Galaxies are fundamental constituents of the Universe. They are groups of approximately 105–1013 stars that are gravitationally bound and take part in the general expansion of the Universe. Galaxies have diameters ranging from 10,000 to 200,000 light-years and they possess widely varying gas and interstellar dust contents. The distances between galaxies are typically 100–1000 times their diameters. They represent 108 overdensities above the mean stellar density of the local Universe. In total, galaxy masses range from 106 to 1014 solar mass (M⊙).

Their component stars vary from ∼0.1 M⊙ brown dwarfs that do not undergo thermonuclear fusion, to rapidly evolving stars of at least ∼50 M⊙ and possibly as massive as 100 M⊙. The stars' evolutionary state ranges from protostars undergoing contraction to begin thermonuclear reactions, to main-sequence dwarf stars that fuse hydrogen to helium in their cores, through to red giant stars with expansive gaseous atmospheres. The stellar end-products are white dwarfs, neutron stars and black holes. The specific evolutionary path of a star is governed by its mass. The most massive stars evolve over millions of years, whilst the lowest mass stars can evolve for billions of years.

At least 300,000 years after the Big Bang start of the Universe, structures that would become galaxies began to condense out of primordial hydrogen and helium.

The following will briefly describe the main telescopes and instruments used to obtain the images presented in the atlas or those used in observations discussed in the text. Detectors are not described. Since there exists significant overlap between telescopes used for submillimeter and radio observations these are considered together.

A1 Gamma ray

kT > 500 keV

The Compton Gamma Ray Observatory (CGRO; Figure A.1) was launched into Earth orbit at 450 km altitude on April 5th, 1991 and re-entered on June 4th, 2000. CGRO contained instruments that could detect radiation with energies from 15 keV to 30 GeV and it was the second of NASA's “Great Observatories”.

These instruments included the Energetic Gamma Ray Experiment Telescope (EGRET) that detected events between 20 MeV and 30 GeV with a positional accuracy of ∼1°. The Imaging Compton Telescope (COMPTEL) covered 1–30 MeV with a positional accuracy of ∼2°. Currently, gamma-ray imaging observations of nearby galaxies are restricted due to a small number of recorded events at poor positional accuracy.

The Fermi Gamma Ray Space Telescope (hereafter Fermi), formerly called the Gamma Ray Large Area Space Telescope or GLAST, was launched on June 11th, 2008 into a 560 km altitude orbit. It is detecting radiation between 8 keV and 300 GeV using the primary instrument, the Large Area Telescope (LAT), and the complementary GLAST Burst Monitor (GBM). The LAT has a large field of view, over 2 steradians (one-fifth of the entire sky), can measure the locations of bright sources to within 1 arcminute and is sensitive to photons from 30 MeV to greater than 300 GeV.

In the early 1930s Karl Jansky detected the center of our Galaxy in radio waves at a wavelength of 14.6 m or a frequency of 20.5 MHz.

Until this time all observations of galaxies had been made in the optical region, which has radiation with wavelengths between 330 and 800 nanometers. Galaxies that can be seen by eye, appearing as faint, cloud-like objects, are the two Magellanic Clouds (named after the Portuguese explorer Ferdinand Magellan) in the southern hemisphere and the Andromeda Galaxy (NGC 224 or Messier 31) in the northern hemisphere. Observers with very good eyesight may also detect NGC 598/M 33, in the constellation Triangulum. We also get a myopic view of nearby stars (about 6000 within a few thousand light-years) and a wide-field view of diffuse light, “the Milky Way”, from distant stars, near the plane of our own Galaxy.

Since the early seventeenth century, telescopes have provided higher resolution observations to fainter brightness levels than possible with the eye. In 2009 the world celebrated the International Year of Astronomy (IYA), as sanctioned by the United Nations and directed by the International Astronomical Union partly to celebrate 400 years of telescopic observations since the initial endeavors of Thomas Harriot and Galileo Galilei.

In terms of effective galaxy research tools, we had to wait until the mid-nineteenth century for large telescopes to be built. Observations by eye, photography and by spectroscopy all played vital roles.

This text is aimed at students and researchers in both astronomy and physics. Spectroscopy links the two disciplines; one as the point of application and the other as the basis. However, it is not only students but also advanced researchers engaged in astronomical observations and analysis who often find themselves rather at a loss to interpret the vast array of spectral information that routinely confronts them. It is not readily feasible to reach all the way back into the fundamentals of spectroscopy, while one is involved in detailed and painstaking analysis of an individual spectrum of a given astrophysical object. At the same time (and from the other end of the spectrum, so to speak) physics graduate students are not often exposed to basic astronomy and astrophysics at a level that they are quite capable of understanding, and, indeed, that they may contribute to if so enabled.

Therefore, we feel the need for a textbook that lays out steps that link the mature field of atomic physics, established and developed for well over a century, to the latest areas of research in astronomy. The challenge is recurring and persistent: high-resolution observations made with great effort and cost require high-precision analytical tools, verified and validated theoretically and experimentally.

Historically, the flow of information has been both ways: astrophysics played a leading role in the development of atomic physics, and as one of the first great applications of quantum physics.

Titan before Cassini–Huygens

With a diameter of 5150 km, Titan is the largest satellite of Saturn. It was discovered in 1655 by Christiaan Huygens. The period of rotation of Titan around the Sun is that of Saturn, 29.5 years. With its obliquity of 27°, Saturn has seasons, each of 7 years' duration, and Titan's seasonality is the same thanks to the close alignment of its pole with Saturn's. In addition, Titan turns around Saturn – with synchronous rotation – within 16 Earth-days, thus Titan's solid surface rotates slowly; however, its atmosphere presents a super-rotation due to strong zonal winds. Titan's mean distance from the Sun is that of Saturn's – about 9.5 astronomical units (AU). This corresponds to a received solar flux at the top of its atmosphere just slightly more than 1% of the flux at the Earth. Moreover, distant from Saturn by about 20 Saturnian radii, Titan is far enough from the giant planet to avoid interactions with the rings, but still close enough to allow its atmosphere to interact with the electrons of the magnetosphere of Saturn, which thus play a role in its chemical evolution, together with the solar photons.

Titan is the only satellite in the Solar System having a dense atmosphere. The presence of its atmosphere was suggested in 1907 by José Comas-Sola based on his observations of the centre-to-limb darkening of Titan's disk.

An elaborate radiative transfer treatment (Chapter 9) is necessary for stellar atmospheres through which radiation escapes the star. But that, in a manner of speaking, is only the visible ‘skin’ of the star, with the remainder of the body opaque to the observer. Radiation transport throughout most of the star is therefore fundamentally different from that through the stellar atmosphere. Since radiation is essentially trapped locally, quite different methods need to be employed to determine the opacity in the interior of the star. However, since there is net outward propagation of radiation from the interior to the surface, it must depend on the variation of temperature and pressure with radius, as in Fig. 10.5.

Perhaps nowhere else is the application of large-scale quantum mechanics to astronomy more valuable than in the computation of astrophysical opacities. Whereas the primary problem to be solved is radiation transport in stellar models, the opacities and atomic parameters needed to calculate them are applicable to a wide variety of problems. One interesting example is that of abundances of elements in stars, including the Sun. Observationally, the composition of the star is inferred from spectral measurements of the atmospheres of stars, i.e. surface abundances, because most of the interior of the star is not amenable to direct observation. However, radiative forces acting on certain elements may affect surface abundances that may be considered abnormal in some stars.

Introduction

Survival of microorganisms in outer space, such as resistant bacterial endospores, is affected by harsh environmental conditions including microgravity, space vacuum leading to desiccation, wide variations in temperature and a strong radiation component of both galactic and solar origins (Nicholson et al., 2000). Solar extraterrestrial UV radiation is mostly deleterious due to its UV component consisting of genotoxic UVC (200 < λ < 280 nm) and more energetic vacuum–UV photons (140 < λ < 200 nm) that are able to ionize biomolecules but exhibit very low penetrating features. The galactic cosmic radiation (CGR) is composed predominantly of high-energy protons (85%), electrons, α-particles and high-charge (Z) and energy (E) nuclei (HZE). In addition, solar particle radiation that mostly consists of protons with very small amounts of α-particles and HZE ions is emitted during solar wind and erratic solar flares (Nicholson et al., 2000; Cucinotta et al., 2008). UVC and UVB photons (280 < λ < 320 nm) are, in the absence of shielding, the main lethal components of space radiation. However, an efficient protection against molecular effects of UV radiations is likely to occur when spores are embedded in micrometeorites according to the scenario that has been proposed for allowing interplanetary or interstellar transfer of microorganisms (Mileikowsky et al., 2000; Nicholson et al., 2000). In contrast, under the latter conditions, protection of microorganisms against the damaging effects of CGR, and more precisely, of highly penetrating HZE particles, is at best very limited.

Introduction

Because Earth is the only place where we are certain that life exists, the characteristics of terrestrial life underpin our search for life elsewhere. In essence, the search for extraterrestrial life begins here on Earth. In the mid-twentieth century, early astrobiologists had recognized this reality and began studying life in remote and extreme environments that could be considered as analogues to places on Mars or elsewhere (e.g. Kooistra et al., 1958; Cameron, 1963; Briot et al., 2004). Early work by NASA and the Jet Propulsion Laboratory included studies of arid-soil microbiology in various locations, including the Atacama Desert and the Antarctic Dry Valleys (Cameron et al., 1966; Cameron, 1969; Horowitz et al., 1969; Cameron et al., 1970). Testing of NASA's earliest life-detection instruments also took place at these and other extreme environments (Levin et al., 1962; Levin and Heim, 1965). In parallel, microbiologists were also studying experimentally the survivability and adaptation of microorganisms isolated from desert soils and exposed to Lunar and Martian conditions in the context of forward contamination of the Moon and Mars, as well as towards the possibility of the existence of extraterrestrial life (e.g. Fulton, 1958; Kooistra et al., 1958; Davis and Fulton, 1959; Packer et al., 1963).

In recent years, Earth-based microbiological research, especially in harsh or extreme environments, has greatly expanded our understanding of the nature and limits of life (e.g. Rothschild and Mancinelli, 2001; Steven et al., 2006; Pikuta et al., 2007; Southam et al., 2007).

Spectral formation depends on a variety of intrinsic atom–photon interactions. In addition, external physical conditions, such as temperature, density and abundances of elements determine the observed spectrum. As described in later chapters, spectral analysis is therefore often complicated and it is difficult to ascertain physical effects individually (and even more so collectively). The main aim of this chapter is to provide a unified picture of basic atomic processes that are naturally inter-related, and may be so considered using state-of-the-art methods in atomic physics. A quantum mechanical treatment needs to take the relevant factors into account. An understanding of these is essential, in order to decide the range and validity of various theoretical approximations employed, and the interpretation of astrophysical observations. From a practical standpoint, it is necessary to determine when and to what extent a given effect or process will affect spectral lines under expected or specified physical conditions.

For example, at low temperatures and densities we may expect only the low-lying atomic levels to be excited, which often give rise to infrared (IR) and optical forbidden emission lines. But the presence of a background ultraviolet (UV) radiation field from massive young stars in star-forming regions of molecular clouds (e.g., the Orion Nebula discussed in Chapter 12), may excite low-lying levels via UV absorption to higher levels and subsequent radiative cascade of emission lines that would appear not only in the UV but also contribute to the intensities of the IR/optical lines.