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.

Gravitational instability

We saw in chapter 4 that our universe contains a hierarchy of structures from planetary systems to super clusters of galaxies. Between these two extremes we have stars, galaxies and groups and clusters of galaxies. Any enquiring mind will be faced with the question: how did these structures come into being?

Is it possible that structures like our galaxy have always existed? The answer is ‘no’ for several reasons. To begin with, stars are shining due to nuclear power which runs out after some time. So it is clearly impossible for any single star to have existed for infinite amount of time. One can, of course, recycle the material for a few generations but eventually even this process will come to an end when all the light elements have been exhausted. So clearly, no galaxy can last for ever. Secondly, we saw in chapter 5 that – at the largest scales – the universe is expanding, with the distance between any two galaxies continuously increasing. This led us to a picture of the early universe with matter existing in a form very different from that which we see today. It follows that the structures like galaxies which we see today could not have existed in the early universe, which was much hotter and denser. They must have formed at some finite time in the past.

How early in the evolution of the universe could these structures have formed? As we shall see, we do not have a definite answer to this question. However, we saw in the last chapter that neutral gaseous systems formed when the universe was about 1000 times smaller.

Stars: Nature's nuclear plants

It is said that a man in the street once asked the scientist Descartes the question: ‘Tell me, wise man, how many stars are there in heaven?’ Descartes apparently replied, ‘Idiot! no one can comprehend the incomprehensible’. Well, Descartes was wrong. We today have a fairly reasonable idea about not only the total number of stars but also many of their properties.

To begin with, it is not really all that difficult to count the number of stars visible to the naked eye. It only takes patience, persistence (and a certain kind of madness!) to do this, and many ancient astronomers have done this counting. There are only about 6000 stars which are visible to the naked eye – a number which is quite small by astronomical standards. The Greek astronomer Hipparchus not only counted but also classified the visible stars based on their brightness. The brightest set (about 20 or so) was called the stars of ‘first magnitude’, the next brightest ones were called ‘second magnitude’, etc. The stars which were barely visible to the naked eye, in this scheme, were the 6th magnitude stars. Typically, stars of second magnitude are about 2½ times fainter than those of first magnitude, stars of third magnitude are 2½ times fainter than those of second magnitude, and so on. This way, the sixth magnitude stars are about 100 times fainter than the brightest stars. With powerful telescopes, we can now see stars which are about 2000 million times fainter than the first magnitude stars, and – of course – count them.

Cosmic inventory

Think of a large ship sailing through the ocean carrying a sack of potatoes in its cargo hold. There is a potato bug, inside one of the potatoes, which is trying to understand the nature of the ocean through which the ship is moving. Sir Arthur Eddington, famous British astronomer, once compared man's search for the mysteries of the universe to the activities of the potato bug in the above example. He might have been right as far as the comparison of dimensions went; but he was completely wrong in spirit. The ‘potato bugs’ – called more respectably astronomers and cosmologists – have definitely learnt a lot about the contents and nature of the Cosmos.

If you glance at the sky on a clear night, you will see a vast collection of glittering stars and – possibly – the Moon and a few planets. Maybe you could also identify some familiar constellations like the Big Bear. This might give you the impression that the universe is made of a collection of stars, spiced with the planets and the Moon. No, far from it; there is a lot more to the universe than meets the naked eye!

Each of the stars you see in the sky is like our Sun, and the collection of all these stars is called the ‘Milky Way’ galaxy. Telescopes reveal that the universe contains millions of such galaxies – each made of a vast number of stars – separated by enormous distances. Other galaxies are so far away that we cannot see them with the naked eye.

Expansion of the universe

The cosmic tour which we undertook in the last chapter familiarized us with the various constituents of the universe from the stars to clusters of galaxies. We saw that the largest clusters have sizes of a few megaparsec and are separated typically by a few tens of megaparsec. When viewed at still larger scales, the universe appears to be quite uniform. For example, if we divide the universe into cubical regions, with a side of 100 Mpc, then each of these cubical boxes will contain roughly the same number of galaxies, clusters, etc. distributed in a similar manner. We can say that the universe is homogeneous when viewed at scales of 100 Mpc or larger. The situation is similar to one's perception of the coastline of a country: when seen at close quarters, the coastline is quite ragged, but if we view it from an airplane, it appears to be smooth. The universe has an inhomogeneous distribution of matter at small scales, but when averaged over large scales, it appears to be quite smooth. By taking into account all the galaxies, clusters, etc. which are inside a sufficiently large cubical box, one can arrive at a mean density of matter in the universe. This density turns out to be about 10−30 gm cm−3.

The matter inside any one of our cubical boxes is affected by various forces. From our discussion in chapter 2 we know that the only two forces which can exert influence over a large range are electromagnetism and gravity. Of these two, electromagnetism can affect only electrically charged particles.

A microscopic tour

The physical conditions which exist in the centre of a star, or in the space between galaxies, could be quite different from the conditions which we come across in our everyday life. To understand the properties of, say, a star or a galaxy, we need to understand the nature and behaviour of matter under different conditions. That is, we need to know the basic constituents of matter and the laws which govern their behaviour.

Consider a solid piece of ice, with which you are quite familiar in everyday life. Ice, like most other solids, has a certain rigidity of shape. This is because a solid is made of atoms – which are the fundamental units of matter – arranged in a regular manner. Such a regular arrangement of atoms is called a ‘crystal lattice’, and one may say that most solids have ‘crystalline’ structure (see figure 2.1). Atoms, of course, are extremely tiny, and they are packed fairly closely in a crystal lattice. Along one centimeter of a solid, there will be about one hundred million atoms in a row. Using the notation introduced in the last chapter, we may say that there are 108 atoms along one centimeter of ice. This means that the typical spacing between atoms in a crystal lattice will be about one part in hundred millionth of a centimeter, i.e., about 1/100 000 000 centimeter. This number is usually written 10−8 cm. The symbol 10−8, with a minus sign before the 8, stands for one part in 108; i.e., one part in 100 000 000.

The broad picture

In the previous chapters, we have explored the conventional thinking of cosmologists and astrophysicists in their attempt to understand the structures in the universe. Some of these attempts have been very successful, while others must be still thought of as theoretical speculations. Since different aspects of structure formation were touched upon in different chapters of this book, it is worthwhile to summarize the conventional picture in a coherent manner.

The key idea behind the models for structure formation lies in treating the formation of small-scale structures like galaxies, clusters, etc. differently from the overall dynamics of the smooth background universe. This is linked to the assumption that, in the past, the universe was very homogeneous with small density fluctuations.

The evolution of the smooth universe is well described by the standard big bang model. Starting from the time when the universe was about one second old, one can follow its evolution till the time when matter and radiation decoupled – which occured when the universe was nearly 400 000 years old. During this epoch, the energies involved in the physical processes ranged from a few million electron volts to a few electron volts. This band of energies has been explored very thoroughly in the laboratory experiments dealing with nuclear physics, atomic physics and condensed matter physics. We understand the physical processes operating at these energy ranges quite well, and it is very unlikely that theoretical models based on this understanding could go wrong. In other words, we can have a reasonable amount of confidence in our description of the universe when it evolved from an age of one second to an age of 400 000 years.

The cosmic rainbow

In chapter 1, we did a rapid survey of the universe, listing its contents, and in chapter 4, we plan to discuss these objects in more detail. You may wonder how such a detailed picture about the universe has been put together. This has been possible because we can now observe the universe in a wide variety of wavebands of the electromagnetic spectrum, and virtually every cosmic object emits radiation in one band or another. In this brief chapter, we shall have a rapid overview of how these observations are made. While describing the observational techniques, we will also mention briefly the astronomical objects which are relevant to these observations. These objects are described in detail in the next chapter, and you could refer back to this cha46 pter after reading chapter 4.

It is rather difficult to ascertain when the first astronomical observation was made. Right from the days of pre-history, human beings have been wondering about the heavens and making note of the phenomena in the skies. The earliest observations, needless to say, were made with the naked eye. With the advent of the optical telescope, one could probe the sky much better and detect objects which were too faint to be seen with the naked eye. As the telescopes improved, the quality of these observations increased.

There is, however, an inherent limitation in these early observations. All these observations were based on visible light. We now know that visible light is an electromagnetic wave whose wavelength is in a particular range.

The subject of cosmology – and our understanding of how structures like galaxies, etc., have formed – have developed considerably in the last two decades or so. Along with this development came an increase in awareness about astronomy and cosmology among the general public, no doubt partly due to the popular press. Given this background, it is certainly desirable to have a book which presents current thinking in the subject of cosmology in a manner understandable to the common reader. This book is intended to provide such a nonmathematical description of this subject to the general reader, at the level of articles in New Scientist or Scientific American. An average reader of these magazines should have no difficulty with this book.

The book is structured as follows: chapter 1 is a gentle introduction to the panorama in our universe, various structures and length scales. Chapter 2 is a rapid overview of the basic physical concepts needed to understand the rest of the book. I have tried to design this chapter in such a manner as to provide the reader with a solid foundation in various concepts, which (s)he will find useful even while reading any other popular article in physical sciences. Chapter 3, I must confess, is a bit of a digression.

Probing the past

The discussion in the last chapter shows that most of the prominent structures in the universe have formed rather recently. In terms of redshifts we may say that galaxy formation probably took place at z < 10. Our understanding of galaxy formation could be vastly improved if we could directly observe structures during their formative phases. Remember that in the case of stars we can directly probe every feature of a stellar life cycle from birth to death; this has helped us to understand stellar evolution quite well. Can we do the same as regards galaxies?

Unfortunately, this task turns out to be very difficult. The life span of a typical star — though large by human standards — is small compared to the age of the universe. This allows one to catch the stars at different stages of their evolution. For galaxies, the timescale is much longer and so we cannot hope to find clear signals for galaxies of different ages. Secondly, the distance scales involved in extragalactic astronomy are enormously large compared to stellar physics. This introduces several observational uncertainties into the study.

In spite of all these difficulties, astronomers have made significant progress in probing the universe during its earlier phases. We saw in chapter 5 that the farther an object is the higher its redshift will be. But since light takes a finite time to travel the distance between a given object and us, what we see today in a distant object is a fossilized record of the past. Consider, for example, a galaxy which is at a distance of one billion light years.

Introduction

The previous chapter contains contributions to the history of science in the field of space plasma physics. It explains how the plasmapause, this peculiar and unexpected magnetospheric frontier, was discovered independently in the late 1950s and early 1960s by two scientists from the two leading countries involved in space exploration. The discoveries were made by using two totally different technical methods of measurement: in situ spacecraft observations and electromagnetic sounding of the magnetosphere. These techniques were both in their infancy at the time.

The main results of electromagnetic sounding of the plasmasphere, from the ground and from satellites, will now be described. In situ satellite particle observations will be outlined in Chapter 3.

In both this chapter and the next the most relevant results will be presented without emphasis upon technical aspects of the experiments. Such aspects are well described in the specialized literature, examples in the case of the whistler method being works by Smith (1961a); Carpenter and Smith (1964); Helliwell (1965); Carpenter and Park (1973); Rycroft (1974a); Y. Corcuff (1975); Tarcsai (1975); P. Corcuff (1977); P. Corcuff, Y. Corcuff and Tarcsai (1977); Park and Carpenter (1978); Bernhardt (1979); Daniell (1986) and Rycroft (1987). An extensive review of the use of whistlers for magnetospheric diagnostics is given by Sazhin, Hayakawa and Bullough (1992).

Initial results

As noted above, Storey (1953) used whistlers for the initial identification of the dense plasmasphere, and Carpenter (1962b) used evidence of unusually low whistler travel times to infer the occurrence of deep, factor-of-∼ 10 depressions in electron density during the severe magnetic storms of the IGY.

Introduction

Preliminary remarks

Besides whistler observations, direct particle measurements from spacecraft (especially from satellites with highly elliptical and geostationary orbits) have contributed significantly to our understanding of the plasmasphere and of its outer boundary, the plasmapause. In particular, such measurements permit us to investigate a number of topics that are not subject to direct observation by radio techniques, including low-energy ion composition, pitch angle distribution, and temperature.

Satellite instruments which have contributed to plasmaspheric studies involve both direct particle flux measurements as well as wave observations. We have already reported in Chapter 2 some results from wave experiments, and in the present chapter will discuss such observations only when they appear to be complementary to direct plasma measurements. Most of our attention will be focused on direct particle measurements, obtained with Langmuir probes, charged particle traps, retarding potential analyzers (RPA), and ion mass spectrometers of different types.

Several problems are inherent in measurements made with the above-mentioned devices. The most serious problems arise when the instruments operate in a very low-density plasma and/or when the energy of the measured particles is very low. Indeed, it is difficult in practice to eliminate all the factors distorting the direct measurements, in spite of the care taken by their developers, including extensive preflight tests and calibration.