This element was discovered in 1923 by D. Coster and G. von Hevesey in Copenhagen, Denmark. The name comes from the Latin name of Copenhagen (Hafnia).

lonization energies

Hfl 6.6 eV, HfII 14.9 eV, HfHI 23.3 eV.

Absorption lines of HfI

The equivalent width of HfI 3616 in the sun is 0.001.

Absorption lines of HfII

The equivalent width of HfII 3505(7) in the sun is 0.024. In one FOIb star, W(3399)=0.008 (Reynolds et al 1988).

Behavior in non-normal stars

Hfl lines are enhanced in some Ba stars, which leads to large overabundances of the order of 0.5–1.0 dex (Lambert 1985, Smith 1984).

Isotopes

Hf has six stable isotopes, 12 short-lived isotopes and isomers and one isotope with a long half life. The stable ones are Hf 174, 176, 177, 178, 179 and 180. In the solar system they represent respectively 0.2%, 5%, 19%, 27%, 14% and 35% of all the Hf. The long-lived isotope Hf182 has a half life of 9x 106 years, which could be useful for radioactive dating.

Origin

Hf isotopes can be produced by several nuclear processes. Hf174 is a pure p process and Hf176 a pure s process product. The other four (Hf 177–180) can be produced by both the r and the s process.

This element is the only one found in the sun (1868, by Janssen) before being detected on earth. It was first isolated by W. Ramsey in London in 1895, and independently by P. T. Cleve and N. A. Langlet in Uppsala. The name comes from the Greek helios (sun). The history of its discovery is given by Frost (1895).

This element was discovered in 1863 in Freiberg, Germany, by F. Reich and H. Richter. The name was chosen because in the spectrum of this element there appears a bright line of indigo color.

lonization energies

InI 5.8 eV, InII 18.9 eV, InIII 28.0 eV.

Absorption lines of InI

The equivalent width of InI 4511 (M.I resonance line) in the sun is 0.004 (Lambert et al. 1969). The In lines are stronger in sunspot spectra. The element is also present in M 2 III stars (Davis 1947).

Emission lines of Inl

InI 4511 appears in emission in long-period variables around maximum light (Merrill 1952, Joy 1954, Deutsch and Merrill 1959).

Behavior in non-normal stars

Cowley et al.(1974) found InII lines in the spectrum of one Ap star of the Cr–Eu–Sr subgroup.

In is enhanced in at least one Ba star (Lambert 1985).

Isotopes

In occurs in two stable forms, In113 and In115. There exist also 32 short-lived isotopes and isomers. In the solar system the frequency of In113 is 4% and that of In115 is 96%.

Origin

In113 is produced by the r, the s or the p process. In115 is only produced by the r process or the s process.

This element was discovered by B. Courtois in Paris in 1811. Its name comes from the Greek iodes (violet).

lonization energies

II 10.4 eV, III 19.1 eV, IIII 33.2 eV.

No I line is observed in the sun. Ionized iodine was found in one Ap star of the Cr-Eu-Sr subgroup by Jaschek and Brandi (1972).

Let us conclude the book with some general reflections.

After reading this book we hope that readers will conclude that our knowledge of the behavior of the chemical elements in stars has advanced significantly since the time of Merrill's (1954) book The Behavior of Chemical Elements in Stars. In this book, which is an excellent summary of the state of knowledge in the early 1950s, Merrill was able to report on 41 elements, whereas in this book we now have some information for 85 elements. In Merrill's time, studies were based upon the optical wavelength region (approximately 3000–8000 Å), but now we are able to provide results from the ultraviolet to the extreme infrared and the radio region. This has helped very much in the detection of molecules – whereas Merrill quoted some 40 molecules, we are able to list 89.

Progress is clearly visible in the production of papers on stellar spectroscopy. In Merrill's time, the annual production was of the order of 100 papers per year, whereas we now find about 1000 papers annually. It is thus legitimate to conclude that knowledge of the subject has advanced significantly.

However, if one regards the subject in more detail, a number of black spots appear on this rosy picture. A perusal of this book shows for instance that, for many of the elements discovered, we still do not know how they behave in stars of the main sequence, not to speak of their behavior in other regions of the HR diagram.

Preliminary Remarks

Petrology is the study of rocks. The prefix, petro, derives from a Greek word meaning rock, or stone. Those unacquainted with geology may be surprised to find that this subdiscipline of geology is huge. All aspects of the study of rocks are included. Not only are various rock types, their mineralogy and chemistry a part of petrology, but also the question of which rocks are where on the surface of the earth, and why. Cosmochemists, of course, must be concerned with lunar rocks and meteorites as well as their terrestrial congeners. Fortunately, lunar and meteoritic petrology are comparatively simple. A fundamental grasp of terrestrial igneous petrology will provide a good introduction to the study of extraterrestrial solids.

The three major divisions of terrestrial rocks are igneous, metamorphic, and sedimentary. Like all classifications these divisions have shortcomings. The boundaries are often ill-defined, and sometimes it is not particularly useful to place a rock in one category rather than another. Igneous rocks are those that have crystallized from a melt or magma. Geologists use the term magma to mean primarily liquid material, which may include some already solidified or not yet melted matter as well as dissolved or imprisoned gases. Lava, on the other hand, may be flowing, or long since frozen magma. Magma becomes lava when it erupts on the surface of the earth.

The signature of an igneous rock is usually evident from a freshly broken surface.

The Identification of Lines in Stellar Spectra

We start our discussion of the analysis of stellar spectra with line identifications. In chemistry, this would be called qualitative analysis. With line identification, we find out which elements and ions are present in stellar atmospheres, leaving the quantitative analysis to subsequent techniques.

The most thoroughly explored stellar spectrum is surely that of the sun. The region from λλ2935–8770, essentially the traditional spectrum available from the ground, is described in a volume by Moore, Minnaert, and Houtgast (1966). This work is a revision of a previous study that was, itself, revised from a still older work. No one at the present time needs to begin the study of any stellar spectrum from first principles. It will be possible in virtually every case to find some at least relevant identification list for a star whose spectrum is similar in nature to the one for which new identifications are desired.

Many of the classical identification studies were done from a list of measured wavelengths, for which the measurer had supplied eye estimates of the line intensities. In modern work one could always have in addition to the list of wavelengths, a tracing of the spectrum, such as the one shown in Figure 12.1. On tracings such as this, one quickly identifies some of the strongest features, such as the hydrogen lines, the resonance lines of Ca II, called H and K, or the strong lines of iron.

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.