The figure that follows provides a short summary of the behavior of the elements.

On the abscissae are given the spectral types and on the ordinates the different elements, specified by the formula (given on the left-hand side) and the atomic number (given on the right-hand side). The continuous line marks those spectral types in which the neutral species is visible, the broken line denotes the presence of the singly ionized species. No higher ionization stages are represented. Blanks indicate elements that are not observed. If a given species of a given element has only been detected in the sun, a black spot is entered in type G. A square in the next to last column indicates that this element has been found to be very strong in some type (or types) of non-normal stars.

To deal with molecules in a book on the behavior of the chemical elements seems a little bit awkward, since the information on molecules seems apparently unsuitable to add to our knowledge of the behavior of the elements. There are, however, several good reasons for dealing with molecules. The first is that, for some elements, we only know their existence through the molecules in which they participate. This is for instance the case for fluorine. In the second place, molecules permit easy study of the different isotopes of an element. As an example consider C, whose isotopes were first studied through the molecular bands of C12/C12 and C12/C13 (Sanford 1940). The study of isotopes is an extremely fruitful subject, which is by no means exhausted. In the third place, molecules are very sensitive indicators of spectral peculiarities. For instance the first indications of the non-uniform composition of stellar atmospheres came through the study of molecules in late type stars. It was found that the division of these stars into two families, one with carbonated molecules (then called R- and N-type stars) and the other group with metallic oxides (the M- and S-type stars) could not be explained by a uniform chemical composition. That such an abundance effect is more easily observed in molecules rather than in atoms or ions, is due to the fact that molecules are very sensitive to differences in chemical composition.

These three arguments largely justify the existence of a separate section on molecules. Another independent reason is that we have been unable to find an updated summary on molecules in stars.

Traditionally the absorption line spectrum of a star has been explained as originating in a (cooler) layer, which surrounds the (hotter) stellar photosphere. The observation of many (unexpected) emission lines from various kinds of objects contradicts such a simple picture. Through the improved observational facilities that have become available in recent decades we now have a much clearer picture of the outer layers of the stars. Since these layers were first observed in the sun, the best introduction is to start with a description of what happens in the sun. This book is certainly not the place for a detailed discussion of the chromosphere and the corona. We shall provide only some information concerning matters that are of interest for the study of stellar spectra. The reader can find more detailed treatment of the solar chromosphere (and also the corona) in many books, like for instance Astrophysics of the Sun by Zirin (1988) and Spectroscopy of Astrophysical Plasmas edited by Dalgarno and Layzer (1987).

The solar chromosphere

The solar chromosphere is the layer enclosing the sun's photosphere. Temperatures diminish in the outer layers of the sun and reach a minimum at what is conventionally called the photosphere. Beyond this layer, temperatures rise again with distance above the sun – first slowly and later very rapidly, until a standstill is reached at about 300000 kelvin. The first part of the atmosphere, where the temperature rise is slow, is called the ‘chromosphere’. The part where temperature rises rapidly is called the ‘transition region’. Grosso modo the chromosphere extends up to 10000 kelvin and the transition region up to 300000 kelvin.

This element was discovered by A. G. Ekeberg at Uppsala, Sweden. The name comes from Tantalus, who was a god, father of Niobe.

lonization energies

Tal 7.9 eV, Tall 15.5 eV.

Behavior in stars

Tal is present in some Ba stars (Lambert 1985, Smith 1984).

Isotopes

Ta has two stable isotopes, Ta180 and Ta181, as well as 16 short-lived isotopes and isomers. In the solar system Ta181 occurs with a frequency of 99.99%.

Origin

Ta180 can be produced by the p process, the r process or the s process, Ta181 only by the r process and the s process.

This element was discovered by C. Perrier and E. Segre in Palermo, Italy, in 1937. Its name alludes to the Greek technikos (artificial).

Ionization energies

TeI 7.3 eV, TeII 15.3 eV, TcIII 29.5 eV.

No Tc lines have been observed in the sun.

Behavior in stars

Tel has its resonance lines at 4238, 4262 and 4297. It was identified by Merrill (1952) in S-type stars. The interest of this element lies in the fact that it is unstable. Its longest lived isotope has a half life of 2.6 × 106 years. The element can be produced by a neutron flux operating on iron-peak elements (Fe, Co, Ni). Such a neutron flux can be either rapid (the r process) or slow (the s process) and the isotopic form in which Tc is formed depends on the time scale in which the neutron flux occurred.

Line identification is one of the basic operations in stellar spectroscopy, which aims at determining the element responsible for each observed spectral line.

The process of identification is based upon

1. (a) measurement of accurate wavelengths of the lines,

2. (b) measurements of the intensity of the lines (or line intensity estimates on a personal scale) and

3. (c) use of line identification tables.

With modern detectors and reduction techniques points (a) and (b) offer no great difficulty, if carried out on adequate observational material – i.e. small plate factors, high resolving power and high signal-to-noise ratio. We should mention an important point concerning the resolving power of the spectrograph, namely that one can only see as separate features two lines that are more separated than the resolving power. The latter is usually fixed by the size of the photographic grain or the pixel size and is normally much larger than the accuracy with which one can measure the position of the barycenter of the (unresolved) line.

On the other hand, point (c) presents a certain number of problems. Very often identifications start with use of the classical paper ‘A multiplet table of astrophysical interest’ (Moore 1945). Such a practice leaves out all laboratory work that has been done since then and that often represents a considerable improvement over Miss Moore's tables. The fact that most of this recent work is not used by astrophysicists is due to the absence of an updated version of wavelength tables, arranged in a similar way to that of Moore. The following compilations are available.

This element was first isolated by H. Moissan in Paris in 1886. Its name comes from the Latin fluere (to flow).

lonization energies

FI 17.4 eV, FII 35.0 eV, FIII 62.7 eV.

F in stars

The presence of some faint FII lines in B-type stars has been announced, but not confirmed. Beals (1951) observed absorption lines of FII in P Cyg.

The best evidence for the presence of F in stars is the molecule HF, which is present in the sun and in K 5 III to M 4III stars (Jorissen et al. 1992).

Behavior in non-normal stars

In S, SC and Cstars the HF feature is enhanced, which leads to an overabundance of F by up to 0.6 dex (Jorissen et al. 1992).

Isotopes

F has only one stable isotope, namely F19, and five short-lived ones.

Origin

F19 is produced by hot hydrogen burning.

This part is divided into 80 sections, each devoted to one element. The elements are ordered by their English names. Cross references between names, formulae and atomic numbers are given in the appendices (part four, tables 2–4).

For the description of the behavior of each element we have kept a constant order, which is as follows.

1. 1 The formula, the English name and the atomic number.

2. 2 A short history of the discovery of the element and an explanation of its name and/or abbreviation. If not indicated otherwise, these data are taken from Elmsley (1989).

3. 3 The ionization energies, in electron-volts (eV), rounded off to the nearest tenth of an electron-volt. The data were taken from Samsonov (1968) and Elmsley (1989). We have only quoted those ionization stages that are observed in stars, plus the next one.

4. 4 We then start a description of the behavior of the lines of each ionization stage (i.e. each species), beginning with the neutral lines. In order to be as precise as possible, we have given the behavior of the equivalent widths of one or more absorption lines with spectral type. Wavelengths and equivalent widths are given in ångström units (1 Å=10−8 cm=0.1 nm), except when stated otherwise. These values are listed in tables and, if sufficient values are available, plotted in figures. The figures provide only smoothed curves, based upon the values listed in the tables. All figures give the spectral types on the abscissae and the equivalent width (in ångström units) on the ordinates.

If data exist in the literature, separate curves have been plotted for dwarfs and supergiants and sometimes also for giants.

The selection of objects included in this book is of considerable importance and we shall make some comments on this point.

To trace the curves of equivalent widths for objects of different spectral types and luminosity classes we have relied only upon MK standard stars. This was done for three reasons. First of all, because MK standards have well-established classifications. Second, they have also been closely scrutinized for spectrum anomalies (both for general line weakening and/or for abnormal line strengths of individual elements). Third, they have been observed so often that spectrum variability can be excluded. If, in the papers from which the equivalent widths were taken, the authors have observed other stars besides the MK standards, then these were not taken into account. When nothing else was available, in a few explicitly mentioned cases, we have accepted stars that are not MK standards.

The use of spectral types and luminosity classes, rather than temperatures and log g values, is justified by the fact that the former are observable characteristics of stellar spectra and are thus stable, whereas temperatures and gravities are fitted parameters and are thus subject to changes in calibration.

The restriction to MK standards has been applied consistently with regard to the behavior of atomic lines. They could, however, not be applied in the chapter on molecules, because the number of stars in which molecules have been measured is too small to allow such stringency in the selection of stars. So we have taken whatever data were available, provided that the classification of the stars was in the MK system.

Conversion of the measured continuum, of the equivalent widths and of the line profiles into physically significant parameters, such as temperature, gravity, turbulence and chemical abundances of elements is the domain of stellar atmosphere studies, for which a number of good textbooks exist, for example

The Observation and Analysis of Stellar Photospheres by D. Gray, second edition, 1992, Cambridge University Press and

Stellar Atmospheres by D. Mihalas, second edition, 1978, Freeman.

Analysis of stellar atmospheres falls outside the scope of this book. Nevertheless, some comments on the precision of the abundances, estimates of abundances and catalogs are necessary.

The precision of the abundances

Abundances are quoted in dex and usually they are given to tenths of a dex (see part two, chapter 2). It is thus often assumed that the precision of the abundances is 0.1 dex. Such a precision is, however, not attainable except in very special cases. A good illustration of the problems that are encountered is given in the proceedings of the workshop on ‘Elemental abundance analysis’ (Adelman and Lanz 1987). At this meeting the same initial data of two sharp-lined Ap stars were provided to several groups of astronomers. The abundances obtained by the different groups show that uncertainties still exist at the level of 0.1–0.2 dex, due to different physical parameters used, differences in computer codes and so on.

If this is true for stars for which the observational data were strictly the same, it is clear that, for stars observed by different observers with different equipment, using different sets of atomic data and different computer codes, precision can only decrease.

This element was discovered by F. Stromeyer in 1817 in Göttingen, Germany. The name comes from the Latin cadmia (calomine, a zinc carbonate). The element was isolated from impurities in calomine.

lonization energies

Cdl 9.0 eV, Cdll 16.9 eV, Cdlll 37.4 eV.

Absorption lines of Cdl

The equivalent width of the Cdl line 5085(2) in the sun is 0.0007.

Behavior in non-normal stars

Jaschek and Brandi (1972) detected Cdl and Cdll in one Ap star of the Cr–Eu– Sr subgroup.

Cdll was detected in two Am stars by Boyartchuk and Snow (1978) with ff(2265)=0.018. Sadakane (1991) found two Cdll lines (2144 and 2265, both of M.I) in the spectrum of another Am star. Isotopes Cd has eight stable isotopes – Cd 106, 108, 110, 111, 112, 113, 114 and 116, as well as 14 short-lived isotopes and isomers. In the solar system one finds respectively Cd 106 (1%), 108 (1%), 110 (12%), 111 (13%), 112 (24%), 113 (12%), 114 (29%) and 116 (7%).

Origin

Cd is produced by several processes: Cd106 and Cd108 by the p, Cd110 by the s, Cd116 by the r and Cd 111, 112, 113 and 114 by either the r or the s process.

This element was first isolated by H. Davy in London in 1808. The name comes from the Latin calx (lime).

lonization energies

Cal 6.1 eV, Call 11.9 eV, Calll 50.9 eV.

Ca is an element of great astrophysical importance. Lines of this element are seen in stars of practically all spectral types, except the hottest ones. If Call lines are visible in O– and B–type stars then they are due to the interstellar medium, or, more infrequently, to circumstellar envelopes.