Introduction

These words describe the enduring relationship between astronomy and atomic physics and, in particular, the wonderful unity of laboratory and astrophysical plasmas. In this chapter we will explore the laboratory part of this unity and will discuss the measurements of some parameters and processes which are important to astronomy and atomic physics.

The determination of chemical abundances and the calculation of model atmospheres for the Sun and stars have become increasingly sophisticated since the landmark work of Goldberg and colleagues. However, despite the impressive progress, calculations based on the best existing atomic data, chemical abundances, and model atmospheres do not reproduce the measured, high-resolution, ultraviolet, solar spectrum nor do they match the center-to-limb variations. When high-resolution, ultraviolet, solar spectra (see Fig. 12.1) are examined it is apparent that the complex of overlapping and nearby lines (‘line blanketing’) could be responsible for much of the discrepancy between observations and calculations.

Overview

A normal main-sequence or red giant star possesses a photosphere, an outer layer of the stellar atmosphere that generates the optical photons that we observe and which is gravitationally bound to the star. Physical conditions in the layer above the photosphere vary greatly among mainsequence stars of different mass and among giant stars, many of which undergo large amplitude photospheric pulsations. Many stars possess chromospheres and, perhaps, also coronae. Mass motions are complicated with some matter falling back onto the photosphere and some being ejected to infinity. It is somewhere in this unsettled region that the outflowing circumstellar shell (CS) begins. For most stars, the circumstellar matter is at least partially transparent and the photospheric spectrum may be observed at wavelengths characteristic of the photospheric temperature. For a few stars with very large mass loss rates, such as IRC + 10216 and CRL 3068, the CS is sufficiently opaque that the photosphere is largely invisible from near infrared to ultraviolet wavelengths. For these stars we see a false photosphere at infrared wavelengths. That is, dust grains that have formed in the outflowing circumstellar gas have sufficient opacity to absorb essentially all of the true photospheric radiation. We then see a cool, roughly blackbody, emission spectrum characteristic of the temperature of dust grains in the inner portions of the CS.

The nature of astrophysical shocks

Energetic phenomena are common in the Universe: Stars inject matter into the interstellar medium via high-velocity winds at their birth and during their lifetime, and massive stars explode as supernovae when they die. Observations of active galaxies reveal rapidly expanding radio sources and jets which appear to be blasting into the ambient medium. On yet larger scales, over-dense regions of the Universe draw matter in at high velocity by gravitational attraction, whereas under-dense regions can expand at high velocity into their surroundings. In many of these cases, the sound speed of the ambient medium is less than the velocity of the gas expanding into it, especially if radiative cooling is efficient. In this case there is no way for the ambient medium to respond to the energy injection in a smooth way. Instead a near discontinuity is produced, a shock, which suddenly accelerates, heats, and compresses the ambient gas. Shocks often govern the dynamics of astrophysical plasmas: they transmit energy from stars to the interstellar medium, they can compress gas past the point of gravitational instability so that stars form, they may terminate the growth of protostars by driving the ambient gas away, they efficiently destroy dust grains and thereby determine interstellar gas phase abundances, and they are efficient accelerators of highenergy particles. Shocks are particularly important in astronomy because they heat gas and cause it to radiate, providing valuable diagnostics for the underlying energetic phenomena and for the ambient medium%.

Introduction

The study of coronal (T≳ 106 K) interstellar gas is a relatively new branch of astronomy. Before the 1970s, there was little direct evidence for such gas, although theoretical models predicted that it should be found in the interiors of supernova shells. In 1956, Spitzer made the prescient suggestion that the galaxy would likely possess a hot corona much like the solar corona. By the early 1970s, a series of rocket experiments had shown that the Milky Way was glowing in soft X-rays, indicating that coronal gas was pervasive in the interstellar medium; this interpretation was supported by observations by the Copernicus satellite of the interstellar absorption line O VI λ1035, showing that this tracer of high-temperature gas was extensively distributed throughout the galaxy.

We now have good maps of the brightness and temperature distribution of the soft X-ray emission from the Milky Way. With X-ray telescopes we have seen emission from coronal gas in elliptical galaxies and between the galaxies in clusters. As a result of these observations, the theory of coronal interstellar gas has advanced rapidly. The atomic processes that determine the local temperature, ionization, and spectral emissivity of the gas have been studied in detail. We have also learned much about the energy sources and macroscopic processes that control the global properties of the interstellar gas. It is now clear that the coronal gas in the Milky Way is produced mainly by the blast waves from supernova explosions, although stellar winds and compact X-ray sources may dominate in specific locales.

Introduction

Quasistellar radio sources were first recognized as objects at very large distances in 1963, when Schmidt identified several nebular emission lines in the spectrum of the stellar appearing, thirteenth magnitude object 3C273, and measured its redshift as z = 0.158. Greenstein and Matthews soon identified similar lines in 3C 48, with V = 16.2, giving z = 0.367. It was immediately clear that these highly luminous objects are beacons that can be observed out to the distant reaches of the universe. For many years OQ 172, with z = 3.53 and V= 17.8, was the most distant object known, until in 1982 an even larger redshift, z = 3.78, was measured for PKS 2000-330, a quasar with red magnitude 17.3. However, observations also quickly showed that all quasars and QSOs do not have the same absolute magnitude; like stars, they are spread over an enormous range in luminosity. Hence we can only hope that ‘understanding’ quasars, through the study of their spectra, will mean not only understanding the strongest energy sources we know in the universe, but also recognizing their absolute magnitudes from their spectra and thus determining their distances.

Long before the first quasars were discovered, galaxies with generally similar emission-line spectra were known. In 1908, Fath, working with a small slitless spectrograph on the Crossley reflector of Lick Observatory, recognized five nebular emission lines (in addition to weak Hβ) in the spectrum of NGC 1068, lines we now know as [O II] λ3727, [Ne III] λ3869, and [O III] λλ4363, 4959, 5007.

The authors and editors of Spectroscopy of Astrophysical Plasmas dedicate this book to Leo Goldberg, who, fifty years ago, recognized both the fundamental role of spectroscopy in the observation and interpretation of astrophysical objects and the essential supporting role of basic laboratory and theoretical studies of atomic and molecular spectra. Leo recognized the importance of all regions of the electromagnetic spectrum from radio waves to gamma rays and he made unique original contributions in ultraviolet, visible, infrared, millimeter and radio astronomy. He fostered the careers of many astronomers, several of whom are authors of this book. His understanding of the value of observations at all wavelengths led him to become a persuasive and influential advocate of space astronomy. At Harvard University he created a research group that was at the forefront of ultraviolet observations, particularly of the Sun.

Atoms, Stars and Nebulae, the title of Leo's first book, written with Lawrence Aller, sums up the main themes of Leo Goldberg's remarkable and still flourishing scientific career. In the early 1930s, when Leo embarked on that career, the ‘new physics’ – quantum mechanics – was still terra incognita for most astronomers. A few, however, had recognized its possibilities years earlier. They saw that quantum mechanics could make possible a quantitative understanding of the structure, composition, and physical conditions of stellar atmospheres and interiors, planetary nebulae, and the interstellar medium; and they set out to do something about it. One of these far-sighted and energetic people was Donald H. Menzel, and Leo Goldberg was a member of the first generation of Menzel's students at Harvard.

Introduction: the nature of the solar chromosphere

The solar chromosphere owes its name to the brilliant red emission seen from the region just above the limb at times of total solar eclipse. The red emission is due to the overwhelming contribution of the Hα hydrogen line at the wavelength 6563 Å. This line emission is produced in large part by scattering of photospheric radiation from hydrogen atoms in the chromosphere, and gives little information about the chromospheric temperature. However, other emission lines such as the D3 line at 5876 Å (the discovery of which gave the name ‘helium’ to the responsible element) indicate a chromospheric temperature considerably higher than the temperature of the underlying photosphere.

Radiative transfer theory, however, indicates that in an atmosphere in radiative equilibrium the temperature generally decreases with height, reaching a surface value near 4300 K in the case of the Sun. Since the temperature in the chromosphere is substantially higher than this value, there must be a source of non-radiative energy to heat it. The height of the chromosphere seen above the solar limb is many times greater than the density scale height appropriate to chromospheric temperatures; this extension and observed rapid motions both indicate that the chromosphere is in a state of intense dynamic activity. The total energy fed into the chromosphere as heat and kinetic energy is about 4 × 106 erg cm-2 s-1, or about 10-4 of the solar luminosity. A small fraction of this large flux of energy penetrates even higher and heats the corona.

Introduction

The outbursts of symbiotic stars have generally attracted much attention, and each eruption is a new adventure in complexity. Since the hot component is the source of the instability, observations acquired throughout the event reveal little of the nature of the cool giant star. Indeed, infrared photometry of CH Cyg and CI Cyg suggests the giant maintains a roughly constant luminosity during a complete outburst cycle (Swings and Allen 1972; Szkody 1977; Kenyon and Gallagher 1983; Luud 1980, and references therein). Thus, the giant appears not to play a direct role in the evolution of an eruption, although it certainly has a good view!!

A typical symbiotic eruption commences with a 2-7 mag rise in visual brightness, which is accompanied by remarkable spectroscopic changes (Belyakina 1979). The B-V color decreases from 1.0-1.5 to 0.0-0.6, suggesting that a relatively hot source dominates the light at visual maximum. While the evolution of B-V as a function of V may be similar for nearly all outbursting symbiotic stars, the duration of and the spectroscopic changes observed in a given outburst naturally divide these systems into three distinct groups:

1. (i) recurrent novae (e.g., T CrB),

2. (ii) classical symbiotic stars (e.g., CI Cyg), and

3. (iii) symbiotic novae (e.g., AG Peg; also known as very slow novae).

Representative light curves for each group are shown in Figures 5.1-5.3; the outbursts of the slow nova-like systems require many decades, while those of CI Cyg and other classical symbiotic systems are much more rapid. The eruptions of recurrent novae are extremely rapid, lasting less than a few months.

Introduction

Initial spectroscopic observations of symbiotic stars revealed them to be peculiar beasts which somehow combine the absorption features of a red giant with the emission lines of a planetary nebula (Figure 1.1). Analysis of their photometric behavior led to additional confusion, as long, apparently quiescent, intervals would be rudely interrupted by 2-3 mag outbursts (Figure 1.2). While the eruptions have captured the imagination of many researchers, a few groups recognized the value of acquiring photometric and spectroscopic data during the lengthy quiescent phases. These historical data, briefly summarized in the Appendix, clarify the long-term behavior of these systems, and provide strong support for the binary models proposed by Berman, Hogg, and Kuiper.

The field of symbiotic stars has grown tremendously in the past twenty years, and data covering a significant portion of the electromagnetic spectrum have been acquired for a number of systems. This information allows definitive tests of the basic models developed in the 1930's and 1940's, and generally supports the basic binary picture discussed in Chapter 2. New results from infrared, optical, and radio photometry demonstrate that symbiotic stars are binary systems (P ∼ 200 days to > 20 years) composed of a late-type giant star and a hot component surrounded by an ionized nebula. Detailed spectroscopic observations have resulted in refined portraits of individual systems, which show a wonderful variety of accreting hot components and mass-losing red giants.

The Discovery of Symbiotic Stars

Most of the stars in the HD catalogue have simple spectra, consisting of a bright continuum and a number of dark absorption lines. These are the “normal” dwarf, giant and supergiant stars that fall in distinct bands in the HR diagram, and form the basis for our current understanding of stellar evolution. However, the spectra of many HD stars display bright emission lines in addition to (or perhaps in place of) the more usual absorption features. Fleming (1912) produced the first comprehensive list of these “stars with peculiar spectra” and grouped them into various categories, including (i) novae, (ii) Otype stars, (iii) stars with bright hydrogen lines, and (iv) long-period variables. A small class of red variables was especially interesting: although their spectral characteristics appeared identical to those of other red long-period variables, their range of variability seemed small for their spectral type. Cannon later isolated another group of red stars with bright lines of H I and He II; these included the irregular variables Z And and SY Mus.

One of the first systematic spectroscopic studies of a peculiar emission-line star was Merrill's (1919) investigation of the enigmatic long-period variable R Aqr. He found conspicuously bright [O III] nebular lines about one month before a predicted maximum. The rest of the spectrum appeared to be that of a normal Md variable, with strong TiO absorption bands and sharp, bright H I lines. Bright emission lines of He II, C III, and N III appeared in 1926 and remained strong until 1933 (Merrill 1936). During these years, the amplitude of the Mira-like brightness variations diminished, although the period of the oscillations remained constant.

Introduction

In the preceding pages, I have tried to present the observed properties of and to develop a theoretical understanding for the group of unusual variables known as symbiotic stars. While excellent reviews of this subject have appeared elsewhere (Allen 1979a; Boyarchuk 1964a, 1975; Feast 1982; Hack 1982; Ilovaisky and Wallerstein 1968; Sahade 1960, 1976), it has been many years since Payne-Gaposchkin (1957) and Swings (1970) discussed the peculiarities of individual symbiotic systems (lists of symbiotic stars also appeared in Bidelman 1954 and Wackerling 1970). Over one hundred stars with combination spectra have been identified since Swings' review; most of these discoveries resulted from detailed objective prism searches conducted by Merrill and Burwell (1933, 1943, 1950), Henize (1967, 1976), and Sanduleak and Stephenson (1973). A number of systems were originally classified as stellar planetary nebulae (e.g., Minkowski 1946, 1948; Haro 1952; Perek 1960, 1962, 1963; Perek and Kohoutek 1967; Frantsman 1962; Thé 1962, 1964; Wray 1966), and later were correctly identified as symbiotic stars by Allen (1979a, b, 1984, and references therein).

In addition to these exciting new discoveries, large amounts of new data have been gathered for the classical symbiotics known by Merrill, especially in the infrared (cf. Allen 1974b, 1980a; Allen and Glass 1974, 1975; Glass and Webster 1973; Feast and Glass 1980; Swings and Allen 1972; Feast, et al. 1983a). It seems appropriate, then, to publish a new summary of individual systems, before the anticipated deluge of material from Space Telescope and other satellite projects buries the recent discoveries made by Einstein, IUE, and various ground-based observatories.