Beat Cepheids continue to pose a number of unanswered questions concerning the behaviour of stars near the short-period end of the Cepheid instability strip; in particular, why is it that such a high proportion of stars with fundamental period ≤ 5 days exhibit simultaneous pulsation in both the fundamental and first-overtone radial modes, and why is it that these stars exhibit such a narrow range of period ratio (P1/P0 ≃ 0.71)? A programme of observation and Fourier analysis of beat Cepheids is presently being undertaken at Mount Stromlo and Siding Spring in an effort to understand these phenomena.

There is now evidence that, for the two best observed stars of this class - U TrA and TU Cas, both the period and the relative mode amplitudes are slowly changing with time. U TrA was observed in detail during the years 1953-59 by Oosterhoff (1957) and Jansen (1962), and a new series of 481 five colour observations of this star was obtained by Faulkner and Shobbrook (1979) in 1977. Using E = Σ(Amp × Freq)2 summed over all the Fourier terms for a given mode as a measure of the pulsational energy associated with that mode, we can use this 20 year timebase to calculate the change in relative mode energy. For U TrA the energy of the overtone is increasing (see Table I). Similarly we can compare the fundamental period derived for these 20 years with an earlier period estimate of Oosterhoff (1957) to obtain the rate of frequency change shown.

Detailed theoretical investigations of stellar pulsation need a mass and a composition before models can be constructed for study. Unfortunately, for the last ten years the masses of the Cepheids, which can be determined with reasonable accuracy by six different methods, have been in dispute. One possible way to resolve these various mass anomalies is to postulate an inhomogeneous composition envelope structure. Thus the principal problem in current theoretical studies is: What are both the masses and envelope compositions of Cepheids?

Stellar evolution theory has indicated that main sequence B stars evolve to become red giants or supergiants and then loop blueward in the Hertzsprung-Russell diagram to cross again the pulsation instability strip the second and third (or even more) times. The second and third crossings are the slowest for models of evolutionary masses between 4 and 13 M⊙. Both below and above this mass the stars become Cepheids (or δ Scuti variables at the lowest masses) on their sole crossing to the red. Massive stars do not become cool enough to become Cepheids if the main sequence mass is above 18 M⊙.

Some twenty years have passed since the decision was taken by the General Assembly of IAU in 1958 to resume a regular control of light variation stability in RR Lyrae stars. With this end in view the revision of elements is carried out at the Odessa Astronomical Observatory yearly, whereas ephemerides of RR Lyrae stars are published at the Krakov Observatory. Now 164 stars brighter than at maximum light visible in the northern hemisphere have entered the list.

Considerable photometric observational material analyzed has shown that some RR Lyrae stars have stable periods of light variation (e.g. TV Lib, UY Cyg, SU Dra), while other ones have periods which vary regularly or proportionally to time (e.g. SV Eri, AV Peg, RV CrB) (Tsesevich, 1966; Firmanyuk, 1976; Romanov et al., 1978). In case of DR And at comparatively small amplitude of O-C variation, about 0.15, a synchronous variation in mean light at maximum nearly by has been found (Romanov, 1978). Unfortunately, necessary observational material is frequently unavailable to obtain the unambiguous classification of period stability character. The property of RR Lyrae stars light variation stability is connected with the weakening of KCall absorption line (Tsesevich, 1966, 1971).

Recently derived bolometric absolute magnitudes are used to investigate the pulsational properties of some classes of red variables. The Mira variables with periods over 200 days have pulsation constants appropriate to overtone pulsatore but the shorter period 135 day Miras seem to be hotter and less luminous and may be fundamental pulsators. The large amplitude Shapley-Nail variables in the SMC are probably fundamental pulsators with masses in the 5-10 solar range. Bright supergiant red variables in the LMC show evidence of the period-luminosity relation expected if they are in the early stages of core helium burning.

The structure of red giant stars allows non-radial oscillation modes which propagate as p-modes near the surface, to propagate below the convection zone as g-modes with very high radial wave number [Dziembowski (1971, 1977), Shibahashi and Osaki (1976)]. Under some conditions the oscillations in these two propagation regions can be treated as virtually independent normal modes [Shibahashi and Osaki (1976)]. This paper examines the situation in which this approximation is not good, and discusses possible observational consequences of the interaction of the two propagation regions.

The linearized differential equations describing non-radial adiabatic oscillations in stars can be written in the form, 1a1b

The invariance of the equations of linear stability of a physical system under a change of sign of the dependent variables implies that we cannot foresee, once an instability is detected, in which direction the system will evolve. In the context of radial dynamical stability this means that the linear analysis does not allow us to discriminate between a subsequent contraction or expansion. To fix up the arrow of the evolution beyond the onset of the instability, a nonlinear analysis is required.

The mass and energy of nova ejecta are 10-5 - 10-4M⊚ and 1044 - 1045 ergs, respectively (Payne-Gaposchkin 1957). Comparison of these quantities with the mass (1 M⊚) and binding energy (1051 erg) of the erupting white dwarf implies that the nova outburst is a surface event. From an average of two or three novae detected each year, it is estimated that the rate of novae is 40-50 per year in our galaxy (Payne-Gaposchkin 1954) . Comparing this rate with a white dwarf birth rate of 2 per year in our galaxy (Weidemann 1968), we conclude that the nova outburst is a recurrent phenomenon (cf. Ford 1978). The recurrent nature also implies that the white dwarf cannot be drastically altered from event to event, thus giving further evidence for a surface event. The argument of recurrency become even stronger when it is realized that observations strongly indicate a close binary structure for the nova candidates—the white dwarf and a red companion (Kraft 1964).

Kraft (1963) proposed the following hypothesis. The red companion overflows its Roche lobe and supplies hydrogen-rich material to an accretion disk around the white dwarf. This material eventually accretes onto the white dwarf, forming an hydrogen-rich envelope whose base is electron-degenerate. As the accretion proceeds, the temperature at the base of this envelope increases. This has little effect on the pressure, but greatly increases the thermonuclear energy generation. The thermonuclear energy generation, in turn, increases the temperature. This positive feedback loop leads to a thermonuclear runaway, which Kraft proposed as the cause of the nova outburst. A large number of theoretical studies based on this model have been carried out. These studies are presented in order of increasing realism and complexity in the following sections.

It is a well accepted idea that Planetary Nebulae ( P.N. ) formation is due to mass ejection from red giant envelopes.

According to the original model, as was proposed by Roxburgh and Lucy in 1967, and established about a year later by Paczynski and Ziolkowski (1968), red giant envelopes become dynamically unstable above a certain boundary luminosity, while their total energy, including the ionization energy, which should be available upon expansion, is positive. This energy is sufficient for ejecting the entire envelope with a velocity similar to those observed in P.N.

In the light of these facts, it was quite natural to look for “one shot” ejection mechanisms in unstable red giant envelopes, using full non-adiabatic dynamical codes. This has been done several times since the adiabatic analysis carried out by Paczynski and Ziolkowski (1968), but with minor success, at least in those cases where natural initial conditions were used.

It was found that as soon as the radial expansion of the envelope becomes comparable to its initial radius, the entire radiative zone, which is located above the partial ionization zone, turns to be almost completely transparent. Thus, most of the released recombination energy is radiated directly out of the star, instead of pushing mass shells outwards. Since this energy leakage is very fast, motion is quickly reversed without resulting in any mass loss. Nevertheless, it was pointed out by Smith and Rose (1972) and later by Wood (1974) that there is a possibility for some mass ejection due to shocks generated within the contracting envelope.

Observed properties of supernovae and of very young supernova remnants provide important clues to answer the question, “which stars become supernovae?” There are three general lines of evidence (l) the statistics of supernovae, (2) the physical state of supernovae near maximum light, and (3) the chemistry of young supernova remnants. These lines of evidence appear to be converging on the view that all the supernova explosions we see, whether Type I or Type II, arise from the destruction of young massive stars.

It is important to note that supernovae fall into two distinct spectroscopic classes: Type I which does not show strong hydrogen lines, and Type II in which hydrogen emission and absorption is prominent (Oke and Searle, 1974). It is possible that the stellar progenitors of these two types are very different, but recent evidence suggests that the similarities in the properties of the two classes are more important.

We show that Rayleigh-Taylor convective overturn of the dynamically formed lepton-trapped core of a supernova is a likely outcome of three sequential events: (1) The bounce or weak reversal shock; (2) the diffusive and convective lepton release from the neutrino-sphere during a fraction of the reversal time (≌ 100 ms); and (3) the rapid (≤ 10 ms) Rayleigh-Taylor growth of the l = 2 mode of an initial rotational perturbation. The overturn releases gravitational energy corresponding to a differential trapped lepton pressure energy of 30 to 50 MeV/nucleon by P dV work in beta equilibrium in a fraction of a millisecond. The resulting kinetic energy of ≌ 7 × 1052 ergs is more than adequate to cause the observed supernova emission. Also, the sudden release of ≌ 7 × 1051 ergs of ˜ 10 MeV neutrinos from the neutrinosphere will cause adequate mass and energy ejection.

Reviews of the mass-loss characteristics of OB stars have been published recently, and the present review therefore emphasizes the A and F stars and very recent results on O and B stars. For the F stars, chromospheric indicators are present in the form of emission lines, seen in visible and ultraviolet wavelengths. Winds are present in A supergiants, but not in main sequence stars, although at least a few of the latter are X-ray sources, indicating the possible existence of coronae. Most OB supergiants are X-ray sources as well, indicating, along with the presence of super-ionization, that these stars have coronae. On the main sequence, the O stars and some B stars (including Be stars in many cases) have mass loss with highly-ionized species in the wind. The winds in the O and B stars are commonly variable. The mass-loss rates do not show a simple dependence on luminosity, contrary to the predictions for radiatively-driven winds.

Until observations in the EUV spectral region became available the existence of hot extended envelopes aroung latetype stars was inferred from the presence of visible and near infrared region lines which cannot be formed in radiative equilibrium within the stellar photospheres. The traditional indicators of stellar chromospheres have been the Ca II H and K lines, particularly the presence of emission components, the He I 10830 A triplet, the He II 4686 A line and emission in various hydrogen Balmer and Paschen lines. When observations from rockets and satellites become possible the strong lines of Mg II at 2800 A and the H Lya line at 1216 A could also be included in the modelling of stellar chromospheres. The new spectra from IUE allow not only stellar chromospheres to be studied but also stellar transition regions, should hot coronae exist.

A number of massive stars of early type is found in X-ray binary systems. The catalog of Bradt et al. (1979) contains 21 sources optically identified with massive stars ranging in spectral type from 06 to B5 out of which 13 are (nearly) unevolved stars and 8 are supergiants. Single stars of this type generally show moderate to strong stellar winds. The X-rays in these binaries originate from accretion onto a compact companion (we restrict the discussion to this type of X-rays).

We consider the compact star as a probe traveling through the stellar wind. This probe enables us to derive useful information about the mass outflow of massive stars.

After presenting the basic data we derive an upper limit to mass loss rates of unevolved early type stars by studying X-ray pulsars. Next we consider theoretical predictions concerning the influence of X-rays on the stellar wind and compare these with the observations. Finally, using new data from IUE, we draw some conclusions about mass loss rates and velocity laws as derived from X-ray binaries.

The solar corona serves as a prototype of the outer atmospheres of all cool stars. Because of its nearness we can study this prototype in more detail than any other example. Considerable progress has been made recently in understanding how the large scale structure of the solar corona controls the genesis of the solar wind and the distribution of slow and fast wind streams throughout the three-dimensional space surrounding the sun. In this review we will discuss some of the progress made in this field during the last few years. We will emphasize the observational data and the inferences that can be made more or less directly from them. T. Holzer will discuss the theoretical aspects of stellar wind acceleration in another paper in this symposium.

The large scale structures of the solar corona consist essentially of three kinds: streamers, active regions and coronal holes. Figure 1 is a familiar photograph of the solar corona, obtained in white light at the total eclipse of 30 June 1973 by the High Altitude Observatory. The streamers are the petal-like structures extending out from the black lunar limb. They taper to narrow radial spikes that have been traced out as far as 10-12 solar radii (Keller, 1979). Daily measurements of the white light corona at the Mauna Loa Observatory (Hundhausen et al. 1979) and the Pic-du-Midi Observatory (Dollfus et al., 1977) since 1965 show that the streamers are fan-shaped structures that may extend 120° in solar longitude. We see them in various perspectives at the solar limb.

Skylab has demonstrated the dominant role of magnetic fields in the solar atmosphere. The solar wind is not a necessary consequence of the pressure imbalance between the hot corona and the interplanetary medium. High speed solar windstreams are originating in coronal holes where coronal density and temperature are less than in the quiet sun. Older models of the solar wind invoke heat conduction from the hot corona as the prime energy source for the solar wind. However in coronal holes the energy supplied by heat conduction is less than in the ordinary sun while the high speed windstreams require an amount of energy which cannot be supplied by conduction alone.

Chromospheres and coronae in stars appear to require vigorous convection zones just below the surface. If we wish to understand how various dynamical instabilities contribute to the mechanical heating that is required to produce chromospheres, then we must be concerned both with fluid motions in the atmosphere and with the nature of their driving below the surface. One cannot really separate these two subjects. In order to emphasize this link, we will raise some basic questions about convective flows in a stellar envelope and of their penetration into the atmosphere. The significant puzzles between what is observed and what can be theoretically explained should serve to indicate some of the issues that need to be pursued. We will concentrate on the Sun in our discussions: the observations here are sufficiently detailed to provide the explicit challenges to theory unavailable in most other stars. However, we will also turn to A-type stars to illustrate a theoretical procedure for describing convection that may do better than the mixing-length approach in predicting the vertical structure in these flows.

The outer atmospheres of stars must be heated by some non-thermal energy flux to produce chromospheres and coronae. We discuss processes which convert the non-thermal energy flux of organized, macroscopic motions into random, microscopic (thermal) motions. Recent advances in our description of the chromosphere velocity field suggest that the acoustic waves observed there transmit very little energy, and hence are probably incapable of heating the upper chromosphere and corona. The apparent failure of this long held mechanism and the growing appreciation of the importance of strong magnetic fields in the chromosphere and corona have led to hypotheses of heating by the dissipation of currents (both oscillatory and quasi-steady). This follows discoveries in laboratory and ionospheric plasmas and work on solar flares, that instabilities can concentrate currents into thin high current density filaments where they dissipate rapidly.

The term stellar wind is used nowadays to describe any more or less continuous mass loss from a star. With the observations made with satellites in recent years it is becoming clear that most stars are undergoing this form of mass loss, though its magnitude can be very different from one star to another. The term stellar wind does not include the more eruptive forms of mass loss such as novae, the ejection of mass in shells or mass loss as a result of flares.

Stellar winds are maintained by energy and momentum deposited in the outer layers of a stellar atmosphere. The deposition of energy causes the heating of a chromosphere and corona, so that the theory of stellar winds cannot really be separated from the theory of coronal heating. Energy and momentum can both be deposited by the same mechanism. For example if a corona is heated by the dissipation of a wave which deposits energy, the same wave can change the momentum of the mean flow through wave pressure and this can happen even in the extreme case of no dissipation of the wave.

The foundation of the theory of stellar winds was laid by Parker (1958) in his theory of the solar wind. A useful review of this work has been given by Parker (1965). The theory of the solar wind in its simplest form is deduced from the equation of motion combined with the equations of continuity and state.

The effects of mass-loss and angular momentum loss on the evolution of massive stars are discussed bringing out the main results as well as the limitations of recent studies. It is pointed out that an acceptable theory of stellar winds in early as well as late type stars is needed as well as a satisfactory assessment of a number of instabilities in these contexts for an adequate understanding of the evolutionary consequences for a wide variety of population I and polulation II stars, which are affected by mass-loss.

In the summer of 1977 the IAU General Secretary requested proposals for Joint Discussions at the XVIIth General Assembly more than two years later. As President of Commission 40 I wrote to other Commission Presidents proposing a J.D. on Extragalactic High Energy Astrophysics. The motivation was as follows, and I quote from my July 22, 1977 letter:

“With the current advances in radio and optical techniques and the tremendous progress to be expected from the satellites HEAO-A and HEAO-B, to be launched respectively in the summer of 1977 and 1978, there should be a good deal of new material on high energy astrophysical phenomena in the extragalactic domain by the time of the 1979 General Assembly. Some of this will be of great cosmological significance and all of it will be astrophysically interesting. It is clear that the X-ray satellite results will get a lot of optical and radio follow-up and it therefore seems appropriate that a Joint Discussion of that topic be arranged at that time.”