The primary reason for studying the dynamics of the local supercluster is for use as a local cosmological test (Silk 1974 Ap. J. 193, 525, Peebles 1976, Ap. J. 205, 318). In addit ion, the measurement of the mass enclosed in the local supercluster provides one of the best measurements of the M/L of galaxy systems on very large scales - a test of the hypothesis that M/L increases with radius for galaxy systems and thus a measure of the unseen matter (neutrino?) content of the universe.

Spherically symmetric models for the dynamic development of a galaxy cluster from an initial overdensity have been carried out numerically, without dissipation or 2-body relaxation but with shell crossings included. The deviation ∆V from pure Hubble Flow of the Local Group, due to the retardation effect of the Virgo cluster and supercluster, has been calculated from a number of different models by Hoffman and Salpeter (Astrophys. J. 263, 1982, in press). The results are somewhat surprising if one takes the point of view of (a) insisting that the dynamic model fit the observed dispersion of galaxy systemic velocities in the core of the Virgo cluster, but (b) allowing the mass to light ratio M/L to be an arbitrary (but smoothly varying) function of distance from the Virgo cluster center. Point (a) essentially fixes the mass density and M/L in the core, but (b) still allows a wide range of values for the cosmological density parameters Ω (proportional to the average M/L far from the Virgo cluster). With this point of view ∆V actually decreases with increasing Ω: If M/L is constant, Ω ≈ 0.3 and ∆V ≈ 250 km s–1 (Hoffman, Olson and Salpeter, Ap. J. 242, 861, 1980); for Ω ~ 0.05, ∆V would exceed 350 km s–1; for Ω = 1, AV could be less than 150 km s–1.

The distribution of rich clusters of galaxies revealed on the photographs taken for the National Geographic Society-Palomar Observatory Sky Survey shows a strong dumpiness on a scale of 100 Mpc (for H = 50 km s–1 Mpc–1), with vast intraclump regions apparently void of rich clusters (Abell 1958; 1961). Studies of the n-point correlation function by Peebles and his associates (e.g., Peebles 1980 and references therein) show that this large-scale dumpiness applies also to individual galaxies. Peebles’ statistical approach does not, of course, indicate individual superclusters or their structures; that kind of information awaited the breakthrough provided by the development of high-speed detectors, with which radial velocities of large numbers of galaxies can be obtained in reasonable observing times.

Recent studies of the spatial distribution of galaxies and of clusters indicate that practically all clusters and a vast majority of galaxies are concentrated into supercusters. The space between superclusters has no rich clusters and very few galaxies. The whole structure is cellular, with cell walls formed from sheetlike superclusters and the empty cell interiors being huge voids.

In this presentation we show how the observed distribution of clusters, and the Local Supercluster in particular, are part of a cellular structure and suggest that the entire pattern of superclustering is cellular. Our study is based on the recently completed CFA redshift survey. All other available redshifts are also used, including cluster redshifts where available.

The astronomer’s requirements for atomic data are so massive that it is not easy to approach this subject in a way that will sound responsible to someone who is not thoroughly familiar with the situation. The simple truth is that there is hardly an atomic species, neutral or ionized, for which additional data is not urgently needed in some domain of astronomy. This comes about partly because of the high temperatures and very long path lengths in astronomical objects which can make them extremely powerful light sources. More generally, we can say that conditions in the astronomer’s violent universe manage to render large volumes of phase space available to the atoms which occupy it.

While the truth is that we need to know “everything,” one task of the present paper is to cast this need in a pragmatic form. This will be done by reviewing certain problems in the very recent history of astronomical spectroscopy, where real progress has come about as a result of new work on atomic spectra. Many of the illustrative problems are taken from the spectroscopy of upper main sequence, chemically peculiar or CP stars.

The data covered are wavelength, line classifications, and energy levels, and the references (numbers in brackets) are mainly supplementary to those supplied by Working Group 4 for the 1982 Report of Commission 14.

New compilations of the energy levels for all the spectra of Fe (Fe I - XXVI) [42] and Co (Co I - XXVII) [43] have been published. Although compilations of the levels for the 235 iron-group spectra (K through Ni) by the NBS AEL Data Center have been published by element during the past five years, the single-volume collection of these levels now under preparation will incorporate fairly extensive additions and revisions for many of the spectra. A compilation of the levels for Si (Si I – XIV) has been completed and should appear early in 1983.

The goal of a quantitative analysis of a stellar spectrum is to derive the physical and chemical state of the stellar atmosphere, i.e. by definition the region emitting the spectrum. Of particular interest are the element abundances, they have to be determined together with the temperatures and densities (pressures) in the atmosphere. A detailed analysis usually is an iterative procedure: a model atmosphere is constructed from reasonable starting values of the parameters (effective temperature, surface gravity, element abundances,...) and used to calculate a theoretical, “synthetic” spectrum. By comparing the observed with the theoretical spectrum, improved stellar parameters are gained for the next iteration step.

Ideally, all atomic data entering the analysis should be known with sufficient accuracy, i.e. errors in the analysis should be due to uncertainties in the assumptions of the models, in the treatment of the velocity fields etc., and not due to insufficient atomic data. In the last decade, the ultraviolet portion of the spectrum below λ ≲ 3200 Å has become accessible by satellites such as Copernicus and the International Ultraviolet Explorer (IUE) with high spectral resolution. Studies of stellar spectra in this new range have revealed the needs for a large amount of atomic data required for the analyses.

Stimulated mostly by the needs of the fusion energy programs, data centres have been established from which information can be obtained on photoionization and recombination processes, electron and ion Impact excitation, ionization and charge transfer cross sections and on electron and proton line-broadening parameters. Regular reports are issued by most of the centres. A partial list follows:

1. (1) Atomic Data for Fusion, Ed. C.F. Barnett, D.H. Crandall and W.L. Wiese: Oak Ridge National Laboratory, P.O. Box X, Bldg, 6003, Oak Ridge, Tenn. 7830, USA.

2. (2) International Bulletin on Atomic and Molecular Data for Fusion, Ed. K. Katsonls: Atomic and Molecular Data Unit, Nuclear Data Section, International Atomic Energy Agency, Wagramerstrasse 5, P.O. Box 100, A-1400 Vienna, Austria.

3. (3) Research Information Center, Institute of Plasma Physics, Nagoya University, Nagoya, Japan, Ed. Y. Itikawa.

4. (4) JAERI Data Center, Division of Plasma Physics, JAERI, 319-11, Tokai-Mura, Japan. K. Ozawa.

In addition, Daresbury Laboratory, Science and Engineering Research Council, Daresbury, Warrington WA4 4AD, England, Eds., W. Eissner and C.J. Noble, publishes an Information Quarterly for Atomic Processes and Applications, which reports continuing developments particularly on calculations of electron impact cross sections.

In order to predict emission line and continuum intensities and assess plasma diagnostics for gaseous nebulae, it is necessary to have accurate atomic parameters. In some Instances the requirements for accuracy are very severe. A variety of data is required:

1. 1) Coefficients of continuous absorption;

2. 2) Direct and dielectronic recombination coefficients;

3. 3) Transition probabilities for permitted lines, so we can accurately predict both recombination line intensities and fluorescent phenomena (e.g., OI λ8446);

4. 4) Transition probabilities for forbidden and intercombination lines;

5. 5) Collision cross sections for levels which give rise to:

   1. a) resonance transitions,

   2. b) intercombination lines,

   3. c) forbidden lines;

6. 6) Charge-exchange cross sections.

In the last decade, research to determine atomic transition probabilities has proceded at a steady pace. Currently, the main needs for transition-probability data arise from astrophysics, magnetic fusion-energy research, and laser development. The principal approaches for the determination of data have undergone further refinement and the major conventional techniques generally contribute data in roughly the same proportions as earlier. However, some significant changes have also occurred: theoretical approaches, especially semi-empirical approximations, have produced vastly increased amounts of data; on the experimental side, determinations of atomic lifetimes by beam-foil spectroscopy have shown a decline, but work utilizing the branching-ratio emission technique as well as lifetime techniques involving laser excitation are on the increase.

This contribution serves to advertise the empirical solar-spectrum determinations of the oscillator strengths of 860 Fe I lines by Gurtovenko and Kostik (1981), by showing that these Kiev data contain just the lines needed in cool-star abundance analyses, and by explaining why they are so good. For details, see Rutten and Kostik (1982) and Rutten and Zwaan (1983).

Figures 1 and 2 show why the Kiev gf-values are so useful. While the precise Oxford measurements (see Simmons and Blackwell, 1982, for references) cover the low-excitation lines (Fig. la) which are mostly too violet and too strong for abundance studies, the Kiev workers have added the best lines in the solar visual, thus precisely the weaker lines of higher excitation one needs (Fig. 1b). Figure 2 demonstrates this for the solar Fe I curve of growth: the Doppler part, which is the only really useful part, is made up of Kiev lines (dots).