The available observational information on the geometrical structure, emission line and continuous spectra, line profiles, radial velocities and linear polarization of Herbig-Haro objects is briefly reviewed. We emphasize the inhomogeneous structure of the “classical” Herbig-Haro objects and the appearance of small “condensations” with radii of ~300-900 a.u. The apparent paradox of the presence of “gaseous nebula type” as well as “reflection type” Herbig-Haro objects is discussed.

A purely empirical model of the regions of line formation is discussed. It shows that regions of low density (Ne ~103 cm-3) cover the space between the condensations and most of the volume of the condensations themselves. Only between 0.1 and 1% of the volume of the condensations is covered by a high density medium (Ne ~ 4 x 104 cm-3, N ~105 cm-3) which, however, contributes very strongly to the formation of the spectrum.

Different theoretical models for the line forming regions are discussed. We strongly favor the shock wave theory in which the emission lines are formed in the cooling regions of (running) shock waves. The general agreement between observations and the new calculations by Raymond is emphasized, and the few remaining discrepancies are discussed. The possibilities of explaining other properties of Herbig-Haro objects (including time scales, sizes and filling factors of condensations) are described.

The Orion population stars are those stars which have only recently formed. It is generally accepted that these stars arose as condensations in the interstellar medium. It follows that they should show strong interactions with that material. In only one case, the nebula near T Tauri (Schwartz 1974) do we have any direct evidence for the interaction between these stars and the interstellar medium. Presumably this is because the typical interactions take place on much larger or much shorter distance scales. The long range interaction of young stars with their environment is discussed in Section II.

It is also generally agreed that Orion population stars are still enshrouded in circumstellar material. There are important questions that have yet to be answered about that material before we can begin to discuss the details of any interaction between these stars and their surroundings. These difficulties are discussed in Section III.

Finally, I have chosen four topics to discuss in somewhat more detail. My aim has been either to urge caution in accepting current interpretations or to put caution aside and suggest new possibilities.

The existence of material in the immediate vicinity of young main sequence stars or even younger objects, protostars, still contracting and evolving towards the main sequence, is related closely to the details of star formation. Although a number of individual idealized problems related to star formation have been well studied, our overall understanding of how interstellar matter evolves into full fledged hydrogen-burning stars can be deemed qualitative at best. The emerging picture is neither complete nor should it be considered the final answer as to how stars form. Let us summarize this picture by a list of “axioms”, valid for star formation (SF) in the plane of our Galaxy.

Simultaneous spectroscopic and photometric observation of the T Tau-like star DR Tau are reported. The spectrograms show a rich emission spectrum. The Balmer lines are strongly variable. They sometimes show a highly complex structure, and inverse P Cygni profiles were observed on two nights. The metallic lines are essentially undisplaced. Simultaneous photometry indicates irregular variations of the blue continuum by as much as 0.5 magnitude on a time-scale of one day and variations of more than two magnitudes on a time-scale of a few days. The observed spectral properties can be explained by assuming that DR Tauri is a protostar of about 1 M⊙ which has not yet reached its hydrostatic equilibrium Hayashi track.

A more complete description of our results is in press in Astronomy and Astrophysics.

During a photographic survey of the Chamaeleon T association in 1971/1972, evidence was found for quasiperiodic light changes of three variable stars, see Mauder and Sosna (1975). The period of 6.2 days for SY Cha is well seen, the periods of 7 days for VZ Cha and of 8 days for TW Cha are less pronounced. Intrinsic variations are present in addition to the cyclic variations. The three stars were classified by Hoffmeister (1963) as T Tauri type stars from their light variations. Objective prism spectra obtained by Henize and Mendoza (1973) confirm this classification, they found conspicuous emission lines. For SY Cha and TW Cha they got slit spectra, too, which show the typical veiling. The stars SY Cha, TW Cha and VZ Cha have been observed in the UBV system from 1974 March 12 until 1974 March 22, using the ESO standard photometer. In Figures 1 - 3 the light and colour curves are given for SY Cha, TH Cha and VZ Cha. Each point is a mean of generally 8 to 12 integrations, each integration lasting 5 seconds.

First results of line profile computations in a collapsing rotating envelope are discussed. As an application, an observed profile of the YY Ori star CoD -35° 10525 is compared to computed profiles. Possible similarities in the formation of YY Ori and T Tau line profiles are briefly mentioned.

A list of 24 T Tauri stars belonging to the YY Orionis subclass is presented. From a statistical analysis it is estimated that at least 75% (and possibly all) UV-excess T Tauri stars are YY Orionis stars. Since about 50% of all known T Tauri stars show a strong UV-excess, the percentage of YY Orionis stars among the T Tauri stars is estimated to be 40 - 50%. This relative high percentage is in good agreement with the present theory of the formation and early evolution of low mass stars.

The YY Orionis stars are a subclass of the T Tauri stars which show in their spectra evidence for infailing envelopes. They are interpreted as protostars in the final stages of their hydrodynamical evolution (Walker 1972, Appenzeller and Wolf 1977, Wolf et al. 1977).

To get more information about the behaviour of these highly variable stars simultaneous spectroscopic and photoelectric (UBV) observations have been carried out in July 1977 at the ESO observatory at La Silla. Such observations of YY Orionis stars have never been done before.

The most frequently observed star during our run was the bright YY Orionis star S CrA. The observational material is still under evaluation and therefore only some preliminary results are presented here. A more extended description and discussion of the observations will be published elsewhere (Mundt and Wolf, 1977).

Three observational studies of YY Ori, the apparently brightest of the YY Ori-type variables in the Orion nebula cluster (Walker 1972), have been carried out. These are:

1. 1. A program of simultaneous spectroscopic and photometric observations. In this program, spectra were obtained on 14 nights at a dispersion of 115 A/mm using a Spectracon image converter mounted at the focus of the Bowen f/1 camera of the coudé spectrograph of the 120-inch reflector (Walker 1966, 1969a). The simultaneous photoelectric observations, on the U B V system, were obtained using the 24-inch Lick reflector.

2. 2. In collaboration with D. Burstein, higher dispersion (0.625 A/channel) observations were obtained on two nights using the Robinson-Wampler scanner at the Cassegrain focus of the 120-inch reflector to investigate in greater detail the structure of the emission and inverse P Cyg absorption features in the spectrum.

3. 3. Spectrophotometry observations were obtained using the instrumentation described in (1) above with a wide slit (4” in width) to investigate the form of the energy distribution in the region from 3200A - 5100A.

230 Å mm-1 spectra of 30 stars of T Tauri or related type were obtained in the 8000-11000 Å region using a grating spectrograph equipped with an S-1 photocathode image-tube. In addition a few T Tauri stars were observed at the same dispersion in the 6200-8800 Å region on IN plate. A preliminary qualitative analysis of the observations leads to possible correlations between the intensities of emissions of the Ca II triplet, Paschen series and He I 10830 Å, and the spectral type, K-L color index of the star, or the [0 I] and Fe II intensities in the visible spectrum.

After the report on flare stars presented by Ambartsumian and the author (1971) at the last Bamberg Colloquium, some success has been achieved in the study of these interesting objects.

Optical observations of stellar flares with high time-resolution undoubtedly proved the explosive character of this phenomenon and the non-thermal nature of the radiation that appears at least at the beginning of the flare. Parallel optical and radio observations of flares in stellar aggregates complicated the interpretation and showed how far we are from an understanding of the physical nature of this phenomenon. Finally, new optical observations of stellar flares in aggregates of different age confirmed the evolutionary significance of this stage in the life of a star and led to the study of certain regularities in the evolution of flare stars.

In a program of systematic search for flare stars in clusters and associations we have carried out a photographic survey in the region of Praesepe. The previously obtained interesting results in the Pleiades cluster and the Orion association stimulated us to continue the work of Haro (1968) and Rosino (1966) in Praesepe. The importance of observations in this cluster is obvious if we consider the following circumstances:

1. 1. Praesepe and Pleiades have nearly the same distance from the Sun, 160 and 125 pc, respectively.

2. 2. There exists a difference in age between the 2 clusters of at least one order of magnitude.

3. 3. The use of Schmidt telescopes allows observations in large cluster regions.

A great number of observational data are desirable for comparison of flare activity in the regions of Pleiades and Praesepe.

U, B, V photometric observations performed in February 1977 at the European Southern Observatory (62 cm Bochum telescope) show that for the star B2IV[e] HD 45677 some of the magnitudes and related color indices may vary by as much as 0.03 mag. over one night, and about 0.1 mag. from night to night. The conclusion derived from our photometric run confirm those published by Feinstein et al (1976), namely that “A variation in V is almost systematically accompanied by a change in B-V but of opposite sign”. We further find that a variation in V is well correlated with a similar change of the U-B index: such a correlation existed in the data of Feinstein and coworkers, although it had not been reported.

Our spectroscopic material, that was obtained simultaneously with the photometry mentioned above, consists, of 38 spectra that cover the blue region at a dispersion of 20 A mm-1. Such an important sample of spectrograms has not yet been analysed in great detail for intensities, radial velocities, etc. . We will only report here on striking variations that were detected in the profiles of the following lines: Hγ, Hδ, Hε, Ca II(H) and Mg II λ 4481.

Continuum fluxes measured with a two-channel spectrum scanner and a 40 A slit are presented, in a differential manner, for six supergiants over the wavelength range 3390-8090 A. Five of these stars presented variability detectable on a time scale of a few days. The stars with the stronger light changes were HD 148379 and HD 160529, both with Hα-emission, whose maximum observed amplitudes amounted to 0.06 and 0.14 mag. The Ha observations done with a 4 A slit gave no indication of strong intensity variations over the same time interval. Some evidence for the presence of a Balmer continuum emission is presented for HD 160529, and we briefly relate it to a possible origin of its Hα-emission profile.

From spectroscopic observations Bessell and Eggen (1972) (hereafter called BE) have detected that HR 5999 (erroneously named HR 6000 by BE) is a shell star of spectral type A7 III-IV with hydrogen lines in emission. Photometric observations show subsequently that HR 5999 varies in brightness up to about 1 mag semiregularly with a period of about 1 month. Eggen (1975) believes that HR 5999 forms a pre-main sequence object.

With HR 6000, the star HR 5999 forms a visual double star system named Δ199; HR 5999 is the southern component. The angular separation is about 45”. BE show that these stars have a common proper motion, and are, therefore, believed to form a physical pair.

The photometric and spectrographic data of the peculiar variable KR Aur have been analyzed. The Hβ line contour shows an accretion with velocity of 3200 km/s. The assumption is made that KR Aur may be a black hole.

The observations of two rather different classical novae, V 1500 Cyg (= Nova Cyg 1975) and NQ Vul (= Nova Vul 1976), are presented and compared. Nova Cyg 1975 is outstanding with respect to absolute magnitude (M = -10), range of brightening (Δm = 19), and speed of evolution (t3 = 3.6 days). Its prenova object had to be fainter than about magnitude 9. The corresponding values for Nova Vul 1976 are rather conservative (M≈-7.5, Δm ≈ 12, t3=40 days, absolute magnitude of the prenova ≈4.5). The light curve of Nova Cyg 1975 is very smooth. Some superimposed photometric variations of a small and slightly variable amplitude of a period of 3.4 hours are most naturally ascribed to a binary nature of Nova Cyg 1975. Nova Vul 1976 has a completely different lighteurve with extremely strong rapid irregular variations of considerably amplitude. There is a remarkable second maximum about 14 days after discovery.

Nova Vulpeculae 1976 was discovered by G.E.D. Alcock (1976) on October 21, 1976 as a star of visual magnitude 6.5 in position (1950):

The visual light curve, plotted together with the color index B-V in fig.1, was obtained using the observations published in the I.A.U. Circulars. It shows a slow fall from maximum, about 2.5 magnitudes in 60 days, with characteristic fluctuations, an abrupt drop of about 3.5 magnitudes in 24 days, a slow recovery of one magnitude in 30 days, and a very gradual decline to minimum. The color index B-V remained almost constant near the mean value +1.1 during the slow early decline.