An Overview

We may divide the meteorites into three broad categories, the stones, the stony-irons, and the irons. Meteorite samples are described as “falls” or “finds.” If a meteorite is observed to fall, and brought to a museum curator, it is called a fall. Finds are meteorites that have not been seen to fall, or at least not by the person who discovers them.

It is now reasonable to speak of two major divisions of meteorites – those discovered in the last several decades in Antarctica, and all of the rest. The Antarctic meteorites have roughly doubled the available samples of solid cosmic debris. It is difficult to know precisely how many independent falls are represented, since all of these samples are finds. However, we must leave differences among the Antarctic and non-Antarctic samples to the references (see Koeberl and Cassidy 1991).

The non-Antarctic meteorites are named by the location of the fall or find. The names are often exotic. For the Antarctic meteorites, locations are also used for the names, but these are supplemented by alphanumerical codes. Most of the world's classified (non-Antarctic) meteorites are listed in the Catalogue of Meteorites (Graham, Bevan, and Hutchison 1985). They include listings for selected Antarctic meteorites.

Most of the meteorites in museums are irons, while the opposite is true of falls – most of the latter are stones.

Introduction

The chemical composition of matter is often the result of factors which, at first, are not at all obvious. Consider the following examples. In certain very stable stellar atmospheres a separation of the chemical elements can take place. One might think that the heavier elements would be the first to sink with respect to the abundant hydrogen, which forms the bulk of (most) stellar matter. In the earth's upper atmosphere, for example, there is a region where the heavy species sink. The number density of a given species roughly follows the law for an isothermal atmosphere, N ∝ exp(–z/H), where z is the altitude, and H the scale height, H = ℜT/gμ. (See the Index for the meaning of symbols not explained in the text.) Thus molecules such as O2 and N2 are concentrated at low altitudes relative to atomic hydrogen and helium.

For the stars in question, the situation is not so simple. There is a competing, upward force due to radiation pressure that can overwhelm gravity. Given time, exotic heavy elements such as mercury or platinum can be pushed up from the envelope of a star and concentrated in the photosphere, where they may be revealed by spectroscopy. In these stars, the abundant, light elements have a tendency to sink! We shall discuss this counterintuitive process in some detail in Chapter 13, since much of the writer's own research has been concerned with its observational consequences.

The Scope of Cosmochemistry

The cosmochemist has two basic tasks. The first is to determine the chemical composition of matter in the material universe. The form of this matter ranges from such mundane materials as terrestrial rocks to distant galaxies. The second task is to understand the reasons for the compositions that are found. While the first of these tasks is basically a matter of analytical chemistry, the second has important evolutionary aspects.

From a logical and historical point of view, cosmochemistry is an extension of the well-established discipline of geochemistry. Victor Goldschmidt, one of the pioneers of modern geochemistry, had a keen interest in abundances of the chemical elements in meteorites and the sun and stars. Goldschmidt (1937) was an early compiler of what became known as the cosmic abundances, in which the analyses of extraterrestrial sources played important roles.

The logical complement to the term geochemistry, in the astronomical domain, would be astrochemistry. This word is frequently in the literature, but typically with a more restricted meaning, such as the formation of molecules in cool interstellar clouds. Solar system astronomers have used terms such as planetary geology or lunar geochemistry rather than astrogeology or astrochemistry.

For many years, astronomers thought of cosmochemistry primarily in terms of the nuclear history of matter, and the search for a standard abundance distribution (SAD). The modern aspects of this work began with Goldschmidt, and were continued by Hans Suess and H. C. Urey (1956).

Introductory Remarks

The philosopher Auguste Comte (1798–1857) asserted that man would never know the chemical composition of the stars. It is therefore ironical that Gustav Kirchhoff (1824–1887) discovered the laws of spectroscopy at about the same time as Comte's death. With the help of the principles articulated by Kirchhoff we now claim knowledge of the composition not only of the nearby stars, but of galaxies so distant that it has taken a substantial fraction of the age of the universe for their light to reach us.

In this chapter we will review the laws of atomic and molecular spectroscopy that enable us to analyze the electromagnetic radiation from space. Naturally, we cannot give a complete account of these rather complicated topics. There is only space to highlight the nomenclature, and in some cases provide heuristic insight into the more important formulae.

The following chapters will deal with the application of atomic and molecular physics to chemical analyses of stars and stellar systems, and interstellar material.

Atomic Spectra: The Nomenclature of LS Coupling

The identification of spectral lines in a star is done with the help of certain reference volumes, the most important of which is possibly C. E. Moore's (1972) A Multiplet Table of Astrophysical Interest. While the basic work appeared as a series of the publications of the Princeton Observatory, the demand for this material was so great that it has gone through one major revision and innumerable reprints and updates.

Introduction

In 1910, the British astronomer Arthur Eddington published an influential monograph with the impressive title of Stellar Movements and the Structure of the Universe. Eddington, who became “The most distinguished astrophysicist of his time” (Chandrasekhar 1983), was only 28 when Stellar Movements was published, but his clarity of exposition and physical insight are readily seen in this small volume. Nevertheless, our present view of the Galaxy in which we live, and the universe around us, is completely different from that limned by Eddington at the end of the century's first decade. Not only were the astronomers of that time uncertain of the nature of the spiral nebulae we now call galaxies (Chapter 16), but they thought the solar system was at the center of our own system of stars.

It had been known since the time of the star gauges (counts) of William Herschel (1738–1822) that faint stars did not increase in number as one would expect, but indicated an “end” of the entire system. Today, at visual wavelengths we can in some sense detect the end of our Galaxy if we look out of the plane, toward its poles. Within the plane, starlight is significantly dimmed by interstellar material – by 1 to 2 magnitudes per kiloparsec (kpc) at visual wavelengths (§§13.3, 14.4).

If we merely count stars as a function of brightness, there is no way to distinguish between the effect of dust and an “end” of the stellar system.

An Introduction to Galactic and Extragalactic Research

In the first half of the twentieth century it became clear that our own stellar system was but one of a very large number of galaxies. To be sure, philosophers such as Immanuel Kant had speculated upon the notion of “island universes,” but it was really the work of Edwin Hubble in the mid 1920's that established the great distance of the Andromeda Nebula (galaxy), and clarified the nature of the extragalactic domain as we know it today. His marvelous book The Realm of the Nebulae (Hubble 1936) is now primarily of historical interest.

Hubble made use of certain highly luminous variable stars known as Cepheids. The Harvard astronomer Henrietta Leavitt had shown that the intrinsic brightnesses of these stars could be obtained from their periods of variation. With this relationship in hand, it was only necessary for Hubble to locate such stars in the Andromeda galaxy, and measure their periods and apparent brightnesses. Their distances followed immediately.

An interesting historical sidelight concerns the errors in the calibration of the absolute brightnesses of the Cepheid variables. By a curious combination of errors, Hubble underestimated the distance to the Andromeda galaxy by a factor between two and three. Harlo Shapley, using similar methods to determine the size of our own system, obtained results that were much more nearly correct because of a cancellation of effects of which he was unaware.

The Concept of Geochemical Classification

The Norwegian geochemist Victor Goldschmidt is the father of the notion of geochemical classifications of the chemical elements. Goldschmidt's (1954) posthumous work Geochemistry is still of great value. His basic aim was to divide the elements into groups which might be identified with the major divisions of the earth during its cooling history. He thought there might be three separate liquid phases, one metal, one silicate, and one primarily iron sulfide. These would be surrounded by a gaseous phase. He classified the elements from their association with, or preference for one or the other of these phases.

Let us begin with a consideration of the chemistry of meteorites and the earth. The earth may be divided into a core, a mantle and a crust. The chemistry of the core must be largely inferred, and this is essentially true for most of the mantle (Ringwood 1975, Pasteris 1984). The upper continental crust is relatively well sampled (Taylor and McLennan 1985), but it is not representative. Meteorites, on the other hand, have been repeatedly and thoroughly analyzed in the laboratory. Moreover, they are thought to be pieces of a broken-up planet, not unlike the earth (see, e.g., McSween 1987). Because of this they have been used to infer the chemistry of the earth as well as of much of the cosmos.

Introduction

The isotopic abundances of cosmic materials may change for a number of reasons. If a substance contains radioactive nuclei, there will be a continual decrease in the parent and a buildup of the daughter isotopes. Bombardment of materials by cosmic rays or other high-energy particles can also alter the isotopic complement of a sample. During radioactive decays or nuclear fission, particles are emitted which can affect the surrounding nuclei. Fission fragments remain in the neighborhood of the parent nuclei. A third possibility is fractionation, by either diffusion or small mass-dependent effects in chemical reactions. All three of these contingencies have been mentioned or intimated previously. We shall now take up certain aspects of these processes in detail.

It will not be possible for us to discuss most of the dating techniques. The interested reader may consult the textbooks of Faure (1986) or Durrance (1986). Richardson and McSween (1989) have an excellent chapter on radioactive dating.

Rubidium–Strontium Dating; Sample and Model Ages

One of the most straightforward methods of age determination makes use of the decay of 87Rb to 87Sr. We shall discuss this particular method here in detail, because of its pedagogical advantages. We shall have time to mention only briefly other methods, some of which are now more actively pursued than rubidium–strontium.

Both rubidium and strontium are geochemically dispersed, that is they occur primarily as impurities in major minerals.

My career as a professional astronomer was some 15 years old when it first became necessary for me to learn something of the new developments in the solar system. At that time, in the mid-1970's, I was about as ignorant of the solar system as one trained in astronomy could possibly be. Worse than that, I had an attitude typical of many astronomers today. Because the field was old, I thought it was dull and uninteresting! Nevertheless, when it became necessary for me to give an introductory course in solar system astronomy, I thought I must try to understand what all the fuss over moon rocks was all about.

Moon rocks are not so terribly different from terrestrial rocks, and so I began to read an introductory geology text. Soon, I was making trips to the building next door to visit the Geology Department. I became an amateur geologist, and a rockhound. On automobile trips I would stop at various rock formations, and bash off samples with a rock hammer. These samples were typically shown to a geologist, sometimes in a nearby university or college, sometimes back at Michigan.

The experience of becoming an amateur geologist was immensely broadening. Not only did I become a great fan of planetary science, but I began to be interested in other areas of astronomy that had never particularly appealed to me. Eventually, I began to realize that there was a single theme behind all of these endeavors – the history of matter.

Introduction

Most stars are converting hydrogen to helium by either the proton–proton chain, or the CNO cycle. We will discuss these processes in §10.2. Nuclear energy is available by the combination or fusion of nuclei less massive than A = 56 (Fe). This nucleus has the maximum binding energy per nucleon. By far the largest fraction of energy in this process is released in the first step, the formation of helium from hydrogen. A rough figure for the ratio of the time that a star spends during hydrogen burning (its main-sequence lifetime) to all other phases of its lifetime is ten to one. The stellar lifetimes are terminated by supernovae explosions, or by quiescent deaths that may involve mass loss followed by the formation of a white dwarf or neutron star.

While the fusion of hydrogen into helium is undoubtedly one phase of nucleosynthesis, it is difficult to know the extent to which the observable helium throughout the universe was made in stars. It is relatively straightforward to demonstrate that the present luminosity of the galaxy, if constant over a lifetime of some 10 billion years, is insufficient to account for the present hydrogen-to-helium ratio. We shall return to this question in §10.7.

Helium, as we shall see shortly, is burned to carbon and oxygen, which can be returned to the interstellar medium in either quiescent mass loss or supernovae explosions. While the helium abundance in the Galaxy is arguably constant, this is surely not the case for carbon.

The General Framework

In this chapter we consider the results of the analytical methods discussed in Chapter 11 insofar as they apply to stars, and to the integrated starlight of some star clusters. Many of these results were obtained with methods that are directly applicable to the analysis of extragalactic systems. We shall postpone a general discussion of the chemical evolution of our own and external galaxies until Chapter 16.

The main pillars of any description of stars and stellar systems are spectral classification, and the Hertzsprung–Russell (H–R) diagram or some variation of it. Classification itself provides mostly information about the temperatures and pressures in the photospheres of stars. Since the majority of stars have rather similar compositions, classification mainly discriminates between the broad categories of normal and peculiar chemistry. The chemically peculiar stars are in many ways the most interesting. The position of a star on the Hertzsprung–Russell diagram (§13.3) indicates its state of evolution. Usually, this tells us something of the chemistry of the star's interior, but occasionally, what has happened in the interior of a star is manifested on its surface. Only a very brief summary of these concepts can be given here.

Most stars belong to double or multiple systems. Double stars were discovered by Sir William Herschel in the late eighteenth century, and a century later, their study was well developed. The marvelous Father Angelo Secchi (1878, p. 228) suggested that perhaps half of the visible stars had physically bound companions.