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.

1. LMC = Large Magellanic Cloud

2. SMC = Small Magellanic Cloud

3. ISM = Interstellar Medium

4. BC = Bolometric Correction

5. Mboi = absolute bolometric magnitude

6. Lbol = bolometric luminosity

7. Teff = effective temperature

8. Γ = –x = Slope of the IMF

9. X,Y,Z = relative content of H, He, and heavy elements (M)

10. [M/H] = logN(M/H)star - logN(M/H)⊙

11. FWHM = Full Width at Half Measure

12. SSPSF = Stochastic Self-Propagating Star Formation

13. LF = Luminosity Function

14. IMF = Initial Mass Function

15. CMD = Colour Magnitude Diagram

16. HR diagram = Hertzsprung Russell Diagram

17. MS = Main Sequence

18. ZAMS = Zero-Age Main Sequence

19. AGB = Asymptotic Giant Branch

20. TP = Thermal Pulse

21. E-AGB = Early AGB

22. TP-AGB = Late AGB

23. RGB = Red-Giant Branch

24. SGB = Sub-Giant Branch

25. HB = Horizontal Branch

26. RGC = Red-Giant Clump

27. PN = Planetary Nebula

28. SN = SuperNova

29. SNR = SuperNova Remnant

As discussed in Chap. 3, the LMC, the SMC and the Galaxy form an interacting system. At times the interactions have had severe effects. Features observed in the LMC and the SMC, including radial velocities of individual objects, are therefore not necessarily determined solely by their rotation and motions.

When the first HI 21 cm-line radial velocities were measured, it was concluded that both Clouds were rotating systems, flattened and tilted and with extensive spiral structures (Kerr and de Vaucouleurs 1955a,b). The rotational motions were at first looked for using the optical centres, 5h24m, −69°.8 for the LMC and 0h51m,−73°. 1 for the SMC. In order to obtain symmetrical rotation curves it was, however, necessary to introduce radio centres of rotation: in the LMC, at 5h20m, −68°.8, appreciably displaced from the optical centre, and in the SMC 1h10m, − 73°.25. The centre of rotation in the LMC has since been a source of much discussion (see Sect. 3.4.4). In the SMC the problems are of another nature.

The structure and kinematics of the LMC

Images of the LMC in most wavelength regions are dominated by radiation from its Extreme Population I constituent (stellar associations, supergiants, etc.) or the connected gas (HII regions, HI complexes, molecular clouds) and dust which display the regions of recent star formation as an asymmetric pattern, not completely at random but with some structure.

Cepheids, RR Lyrae stars, Mira variables, clusters and OB stars have dominated in the attempts to determine accurate distances to the Magellanic Clouds in recent years. The subject is of vital importance not only for Magellanic Cloud research but also for the determination of extragalactic distances (see Feast 1988, van den Bergh 1989, de Vaucouleurs 1993). With the accurate photometry possible nowadays the uncertainties in the distance determinations arise mainly through their dependence on the metallicities and ages of the objects involved and on the interstellar absorption in the line of sight, be it in the Galaxy or in the Clouds themselves.

The reddening, foreground and internal

The interstellar reddening affects virtually all studies of the Clouds. It is frequently difficult to determine its magnitude from the material available and mean values found in the literature are then, by necessity, applied. This may occasionally lead to erroneous results, in particular as the objects used to determine the mean may not be representative of the whole galaxy. In the following, some investigations will be referred to either as giving general ideas about the mean reddening or as examples of typical deviations.

The galactic foreground colour excess towards the Magellanic Clouds has been investigated by Schwering and Israel (1991) on a scale of 48 arcmin from galactic HI maps, using EB-v = 0.17 10−21 N(HI) H cm−2.