Most spiral galaxies are, to a good approximation, oblate spheroids that can be arranged on the sequence Sa–Sb–Sc. However, a significant minority exhibit bar-like structures and may be placed on the SBa–SBb–SBc tine of Hubble's ‘tuning fork’ diagram. Objects, with less pronounced bars, can be arranged on the intermediate sequence S(B)a–S(B)b–S(B)c between normal and barred spirals. Some disks appear to be globally oval (Kormendy 1982, p. 135). Such oval disks are of interest because, like bars, they represent non-axisymmetric distortions of the gravitational potential. An excellent review on the dynamics of barred galaxies has been given by Sellwood & Wilkinson (1993). Infrared imaging shows that some galaxies contain small IR bars (Frogel, Quillen & Pogge 1996). When such small inner bars are found in galaxies with large outer bars there is no correlation between the position angles of the large and small bars.

Bars are important to the dynamical evolution of galaxies (Kormendy 1982, 1993) because they can (a) lose angular momentum to cold disks and dark halos, and (b) gain angular momentum from rapidly rotating bulges. Furthermore, transfer of significant amounts of gas by bars to the nuclear regions of galaxies will increase the central galactic mass concentration, which in turn will make the galactic disk less prone to the development of bar-like distortions. In other words stellar bars can self-destruct (or transform themselves into lenses) by transporting too much gas to the nuclear regions of their parent galaxies.

Yogi Berra once said that ‘you can observe a lot by just watching’. The truth of this aphorism struck me when the first prints of the Palomar Sky Survey started to arrive at the Universitäts Sternwarte in Göttingen, where I was a graduate student in 1955. Just looking at this marvellous atlas immediately showed a number of interesting things that had not been so obvious on the smaller, and less homogeneous, databases that had previously been available: (1) The most luminous galaxies in clusters are ‘pretty’ because they have long well-defined spiral arms, whereas ‘ugly’ spirals of lower luminosity tend to exhibit short patchy arms. (2) Intrinsically faint galaxies generally have lower surface brightnesses than do luminous ones. (3) Galaxies in rich clusters sometimes exhibit peculiarities, like fuzzy spiral arms, that are rare among isolated field galaxies.

In the present volume, which is based on a series of lectures given at the University of Victoria in early 1997, I have tried to provide an up-to-date summary of current ideas on the morphology [morphe = shape] and classification of galaxies. I am indebted to Roberto Abraham for suggesting that I write this review. I also thank Ralf Bender, Scott Tremaine and Stephen Zepf for discussions on the interpretation of the classification of elliptical galaxies, and Guy Worthey and Masafumi Noguchi for discussion of the abundance ratio of elements to iron in normal and barred spirals.

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.