Introduction

Dwarf spheroidals are the most common type of galaxy in the Universe. The fact that they were not discovered until 1938 is entirely due to their feeble luminosity and low surface brightness. Of the 29 galaxies that are known to be located within 1.0 Mpc, approximately half are dwarf spheroidals (dSph). A listing of these Local Group dSph galaxies is given in Table 14. For the sake of completeness the dSph/dE galaxies NGC 147 and NGC 185, which are both brighter than Mv= −15.0, have been included in the table. Since most of the faintest known Local Group members are dwarf spheroidals it is almost certain that additional very faint dSph galaxies remain to be discovered in the Local Group. In particular it seems probable that more dSph companions to M31 will eventually be found. Only three such objects (And I, And II and And III) are presently known (van den Bergh 1972), whereas seven dSph companions (Sgr, UMi, Dra, Scl, Sex, Car, For) are known to be located within 150 kpc of the Galaxy – even though the Milky Way system is less luminous than the Andromeda nebula. It is, of course, possible that the small number of M31 dSph satellites is due to the fact that some dwarf companions to M31 were destroyed by tidal interactions with M32 and NGC 205. For reviews on dwarf spheroidal galaxies the reader is referred to Da Costa (1992), Gallagher & Wyse (1994) and Ferguson & Binggeli (1994).

A short history of stellar astrophysics

Since the Sun is a star it is probably correct to say that stellar astrophysics began with Newton's well-known explanation for the Keplerian laws of planetary motion. Although J. Goodricke observed the eclipsing binary variable Algol (β Persei) in 1782, it was not until 1803 that Sir William Herschel's observations of Castor proved that two stars revolve around each other owing to their mutual gravitational attraction.

The first measurements of stellar parallax were made by F. W. Bessel and F. G. W. Struve in 1838. F. Schlesinger revolutionized stellar distance determinations in 1903 when he introduced photographic parallaxes and thereby enabled astronomers to measure parallaxes to an accuracy of about 0.01 arc seconds. K. Schwarzschild initiated photographic photometry during the years 1904–8. Photoelectric photometry of stars began shortly after the photocell was invented in 1911.

J. Fraunhofer discovered Fraunhofer absorption lines in the solar spectrum in 1814 and subsequently observed similar lines in other stars. In 1860 Kirchhoff formulated the relationship between radiative absorption and emission of radiation which is known as Kirchhoff's law. The Doppler effect and Kirchhoff's law formed the conceptual basis of early studies of stellar atmospheres. The quantum theory of blackbody radiation was introduced by M. Planck in 1900. To a first approximation most stars radiate as blackbodies with superimposed absorption and emission lines. The modern theory of radiative transfer in stellar atmospheres was initiated in 1906 by K. Schwarzschild.

Remarkable progress in understanding stellar phenomena has occurred in recent decades. This textbook discusses in some detail those equations and physical processes that are of greatest relevance to stellar interiors and atmospheres and closely related astrophysics. Motivation for writing this book came from my own research interests and also from teaching graduate astrophysics courses, especially a course on stellar interiors at the University of Maryland. Although the text emphasizes physical principles, astronomical results and unresolved issues are also described.

Introductory material on the history of stellar astrophysics, astronomical observations, star formation and stellar evolution are given in Chapter 1, which also contains a discussion of spectroscopic binaries. Differences between single and binary star evolution have explained a number of interesting observations that are described further in later chapters.

Stellar interiors is one of the most fundamental subjects in astrophysics. Although complicated physical processes are decisive in explaining some predictions of stellar model calculations, the basic principles of stellar interiors do not require a comprehensive knowledge of them. Chapter 2 gives an introductory discussion of the physics and equations of stellar interiors. It also includes a short description of numerical methods.

Statistical physics provides the theoretical basis for much of stellar astrophysics. In Chapter 3 those aspects of statistical physics that are of greatest relevance are developed in some detail. Stellar opacities play a vital role in interpreting observations. Absorption processes are described in Chapter 4.

Morgan (1958) has said that ‘The value of a system of classification depends on its usefulness.’ Using this criterion the Hubble classification system has proved to be of outstanding value because it has provided deep insights into the relationships between galaxy morphology, galactic evolution and stellar populations. However, some classification parameters, such as the r and s varieties in the de Vaucouleurs system, have not yet been tied as firmly to physically significant differences between galaxies (cf. Kormendy (1982)). Furthermore, it is not yet clear if the dichotomy between ordinary and barred spirals allows one to draw any useful conclusions about the past evolutionary history of a particular galaxy.

The Hubble system was designed to provide a framework for the classification of galaxies in nearby regions of the Universe. It is therefore not surprising that it does not provide a useful reference frame for the classification of very distant galaxies (which are viewed at large look-back times), or for galaxies in unusual environments such as the cores of rich clusters. Furthermore, the existence of some classes of objects, such as (1) amorphous/Ir II galaxies, (2) anemic galaxies and (3) cD galaxies, which cannot be ‘shoehorned’ into the Hubble system, suggests that such galaxies have had an unusual evolutionary history. It has also become clear that the Hubble system, which is defined in terms of supergiant prototypes, does not provide a very useful framework for the classification of low-luminosity galaxies.

Pulsational instability

In Chapter 1 we discussed some of the observational properties of periodic variable stars. The instability that drives pulsations in RR Lyrae variables, Cepheids and long-period variables is associated with hydrogen and helium ionization zones. The large heat capacity of these ionization zones causes the phase of maximum luminosity to be delayed by approximately 90° as compared to the phase of minimum radius. Thermonuclear reactions can also cause stars to become pulsationally unstable. Very massive stars and white dwarfs in which thermonuclear runaways are caused by mass accretion from a binary companion become pulsationally unstable as the result of their hydrogen-burning sources. To determine whether a particular star is pulsationally unstable one first determines the structure of the star (i.e. r = r(Mr), P = P(Mr), ρ = ρ(Mr), Lr = Lr(Mr)) and then solves the linearized equation of motion for the oscillatory modes. It is usually adequate to assume that stellar oscillations are adiabatic. If the oscillatory modes of a star have been determined we can evaluate a stability integral which will be derived below. The sign of this stability integral determines whether a particular stellar model is unstable to self-excited oscillations at a particular frequency (eigenmode). We are usually interested only in radial modes of oscillation and in most circumstances only the longest period mode is pulsationally unstable. In β Canis Majoris stars (also known as β Cepheid variables) nonradial oscillatory modes can also become excited.

Stellar clusters occupy a central position in research aimed at the structure and the evolution of our Galaxy and of those of our neighbours in which clusters can be identified. Often the integrated cluster properties, magnitudes, colours, spectra, are the only ones within reach. In the Magellanic Clouds most of the clusters can be sufficiently resolved for the investigation of individual members by photometry and spectroscopy even if the stars in the cores in some cases are too crowded for ground-based observations. As the clusters have a range of age that covers the whole lifetime of the Clouds, this should permit the study of the complete evolution of the Clouds. En route, a number of steps have to be taken. It is necessary to determine their distances, ages, and metallicities, and, before these, their reddening. The latter is difficult to determine for an individual cluster without knowledge of its physical properties, and is, therefore, frequently assumed known. As the reddening is small over most of the Clouds (see Chap. 2), the astronomer may feel entitled to use any low value recommended in one survey or another. However, even a small error in the colour excess, EB-V, may have noticeable effects on the other quantities. Also the distance to the cluster, i.e. to a particular part of the SMC or the LMC, is frequently assumed known or determined by isochrone fittings: isochrones for different compositions and ages are fitted to the main sequences (MSs) and/or the red-giant branches (RGBs) in the colour–magnitude diagrams (CMDs) and the best fitting one is accepted as defining the cluster properties.

The concentrations of luminous, blue stars in the Magellanic Clouds have attracted much attention. Shapley (1956) noted that the large gaseous nebulae in the LMC are frequently associated with groups of stars but also that some of the larger star groups are free of conspicuous nebulosity. As these stellar aggregations were too large to be called clusters or associations in the sense used in our Galaxy he called them ‘Constellations’. He estimated their diameters to be between 250 and 600 pc and their content of blue supergiants, with a red magnitude brighter than 14.0, to be between 14 and 32. A few red supergiant stars were seen in each of them. In the region of 30 Doradus, Shapley (1955) identified a number of red stars by comparing blue (B) and infrared (I) plates and concluded that the very red stars were of spectral class M0 or later. Only two of the 21 most luminous of these stars are in the vicinity of the core of 30 Doradus.

In the SMC Shapley found only one object rich enough to be called a constellation in the sense used for the LMC. It is the aggregate comprising NGC 456, NGC460 and NGC 465 in the Wing area.

Shapley's designation is still used to identify the five most conspicuous stellar aggregates in the LMC. Improved techniques have extended and redefined them and led to the identification of more such formations.

The Magellanic Clouds have been known for thousands of years to the inhabitants of the southern hemisphere. The natives on the South Sea Islands called them the Upper and Lower Clouds of Mist. The Australian aborigines, who referred to the Milky Way as a river or track along which the spirits travelled to the sky-world, considered the Magellanic Clouds as two great black men who sometimes came down to the earth and choked people while they were asleep (McCarthy 1956, p.130). Al Sufi, in his description of the stellar constellations from the 10th century, told about a strange object, A1 Bakr, the White Ox, which is now identified as the Large Magellanic Cloud. Many of the mariners of the Middle Ages noticed the two Clouds. The Italian Corsali described them: ‘We saw two clouds of significant size which move regularly around the pole in a circular course, sometimes going up and sometimes down, with a star midway between them at a distance of 11 degrees from the pole and participating in their movements.’

The two objects were called the Cape Clouds for hundreds of years; they were the most striking objects appearing in the sky when ships approached the Cape of Good Hope. They were of importance for the navigators of that time for localizing the South Pole, where there is no star corresponding to Polaris in the North (see Allen 1980).

It is essential for our understanding of the evolution of the Magellanic System, comprising the LMC and the SMC, the InterCloud (IC) or Bridge region and the Magellanic Stream, to know the motions of its members in the past. The Clouds have a common envelope of neutral hydrogen. This indicates that they have been bound to each other for a long time. It is generally assumed, but not definitely proven, that the Clouds have also been bound to our Galaxy for at least the last 7 Gyr. Most models assume that the Clouds lead the Magellanic Stream. The Magellanic System moves in the gravitational potential of our Galaxy and in the plane defined by the Local Group It is also exposed to ram pressure through its movement in the galactic halo. The influence of our Galaxy ought to be noticeable in the present structure and kinematics of the System.

The interaction between the Clouds has influenced their structure and kinematics severely. It should be possible to trace the effects as pronounced disturbances in the motions of their stellar and gaseous components. Recent astrometric contributions in this field show great promise for the future if still higher accuracy can be achieved. It should be kept in mind in all analyses that results of interactions may be expected everywhere.

Interest in the Magellanic Clouds has grown tremendously over the past four decades. During this period they have been exposed to investigations, interpretations, and speculations with regard to their origin, evolution, structure and content. At times, they have been viewed as more spectacular than they perhaps really are, e.g. suggested to have supermassive stars and peculiar structures; at other times they have been wished far away. Shapley once said (in Galaxies, Harvard University Press, most recent edition 1973, ed. P.W. Hodge) that ‘The Astronomy of galaxies would probably have been ahead by a generation, perhaps by 50 years, if Chance, or Fate, or whatever it is that fixes things as they are had put a typical spiral and a typical elliptical galaxy in the positions now occupied by the Large and Small Magellanic Clouds…. But we must make the best of what we have, and it will soon appear that the best is indeed good. It's marvelous.’ This has indeed been shown to be true. The two irregulars, which differ in so many aspects from our Galaxy, have in particular shown their value as two excellent astrophysical laboratories near at hand.

The study of the Magellanic Clouds has in many ways become more ‘galactic’ than ‘extragalactic’. It is therefore equally impossible to cover all Magellanic Cloud research in detail in one monograph as it would be for our Galaxy.

X-ray emission from the Magellanic Clouds was first observed in a five-minute rocket flight from Johnston Atoll in the South Pacific on October 29, 1968. The LMC was detected as an ∼ 4 σ excess in two adjacent 5° bins, the flux was ∼ 1.5 × 10−9 erg cm−2 s−1 and the spectrum was slightly softer than that of the diffuse background (Mark et al. 1969). Two years later two source regions in the LMC were identified by Price et al. (1971) and emission from the SMC was recorded. The same year Leong et al. (1971) showed that the LMC emission could be resolved by the collimated detector system of the Uhuru satellite into three steady and one possible highly variable source; they were designated LMC X-1, X-2, X-3, and X-4. There was also a possible diffuse emission extending over much of the Cloud. The SMC emission was located to a single, highly variable source, called SMC X-1, in the Wing region. It was the first stellar X-ray source to be confirmed in an external galaxy.

Confirmation of the existence of LMC X-4 was presented in the second Uhuru catalogue (Giaconi et al. 1972).

The LMC sources X-1, X-2, and X-3 were confirmed by Copernicus satellite observations (Rapley and Tuohy 1974) and given more accurate positions. Also Markert and Clark (1975), using the OSO 7 satellite confirmed these three sources, defined an upper limit for LMC X-4, and introduced a fifth source, LMCX-5.