The changes in stellar polarization with time were first discovered in 1958 by GRIGORYAN (1958) for the red supergiant µ Cephei. In subsequent year BEHR (1959) gave a hint in a footnote to his catalogue that the emission-line B-type star α Cas seems to have variable polarization. However, several years had to pass before these discoveries were confirmed and the changes in polarization with time for the red giants and supergiants and for the emission-line B-type stars were firmly established, first by the observations made by SHAKHOVSKOJ (1963, 1964) and then by COYNE, GEHRELS (1967) and the writer (SERKOWSKI 1966, 1968). The changes in polarization with time were also found for the RV Tauri and T Tauri stars, R Coronae Borealis stars, novae, Wolf-Rayet stars, and white dwarfs. The rapid changes in polarization with time indicate that the polarization is not produced in the interstellar medium but in the vicinity of a star; such polarization is called intrinsic.

We know several infrared objects showing large intrinsic polarization (FORBES 1967, APPENZELLER and O’DELL 1967, KRUSZEWSKI 1968, SERKOWSKI 1969, SHAWL 1969, SHAWL and ZELLNER 1970, KRUSZEWSKI 1971, DYCK, FORBES and SHAWL 1971).

In the following discussion I shall speak about the broad category of infrared stars which consists of reddest and coolest giants (M8 or later) and supergiants (M5 or later). This restriction is made because the yellow supergiants seem to have different polarization properties and these properties most likely extend to supergiants as late as M 2 or even M 4. Also it is not clear if normal red variables can be treated together with infrared objects.

The time dependent polarization of the light of some red long period variables can be explained using the simple model of the oblique rotator. Near the magnetic poles the dust, particles in the outer layers of the stellar atmosphere could be aligned in the magnetic field by the DAVIS-GREENSTEIN mechanism. For the stars VCVn, L2 Pup, R Leo and R Peg the spatial orientation of the rotational axis, the angle between the magnetic poles and the rotational axis and the period of the revolution of the magnetic poles around the rotational axis were derived from the observational material. The proposed hypothesis gives an explanation of the essential observational facts concerning the intrinsic polarization of some of the red variables.

Assuming a stellar magnetic dipole-field which aligns particles in the stellar atmosphere the observed polarization of the light of µ Cephei can be explained. The spatial orientation of the rotational axis, the angle between the magnetic axis and the rotational axis and the period of rotation were derived from the observational values. The period of rotation amounts to about 1650d ± 100d. In case aligned graphite flakes are present in the outer layers of the stellar atmosphere the calculated degree of polarization amounts to about 1.5 %. According to the chemical composition of the atmosphere it is more probably that siliciumcarbide particles cause the polarization. In this case the degree of polarization amounts to approximately 0.75%. The portion of polarization originating in interstellar space was estimated to Pi ≈0.8% and the position angle is about Θi ≈45°.

„An astronomer is someone to whom a factor of 10 is important, but a factor of 2 is not.“ If we agree about this, we may say that stellar polarization is 1% and that we wish to measure that polarization to l% accuracy. Polarimeters therefore need an accuracy of 0,01 % (of the total intensity).

In general, stellar polarimetry has been rather a delicate measurement. This is because its nature is differential, i. e. the information required is the small difference between two large observables. There are two general classes of polarimeter, single-detector and two-detector. In the single-detector polarimeters the observables are measured one after the other by the same detector and time-varying attenuations cause errors. In two-detector polarimeters, differences between the two channels cause errors. In both cases extreme stability is required. The solution is to use the single-detector type, but to alternate between the two observables faster than any of the time-varying attenuations. This type of polarimeter uses some sort of polarization modulation.

Stellar polarization between 3000 and 7500 Å has been measured in NGC 2264. Out of 37 stars (including many variables) 4 are definitely polarized, while 3 others show weak polarization. The mean (interstellar) cluster polarization is found to be very small. The highest polarization (4,6 %) is found in a T Tauri star. W 90, a variable below the main sequence, has the highest polarization found among the A and F stars (2, 6%). With the exception of W46, polarization peaks in the blue implying small particle sizes. A good correlation is found between the size of polarization and (other) shell characteristics. These measurements confirm the recently proposed shells in some pre-main sequence stars which appear to be too faint for their color, but fail to detect any new stars with extensive shells. This suggests that some of the scatter in the color-magnitude diagram amongst the A and F stars is real because of an age spread between 1 and 3 X 106 years.

During the past several years, the radio astronomical detection of a large number of complex organic and inorganic molecules has led to an entirely new concept of the chemical composition, kinematics and excitation of the interstellar medium. The microwave molecular transitions detected thus far fall into two natural categories: 1. „normal“ (i.e. nearly LTE) emission and absorption lines, and 2. anomalous or „masered“ emission lines. The first category includes molecules found in both high and low excitation regions whose intensity, linewidth and other characteristics may be described in terms of a physically realistic excitation temperature. Nearly all of the known microwave lines fall into this category. Typically, molecular lines of this type are found in HII regions, dark nebulae and, occasionally, extended circum-stellar shells. Only two molecules populate the second category: OH and H20. Emission from OH lambda doublet transitions in several rotational states and from the 618—523 rotational transition of H20 exhibits many strange characteristics: tremendous intensity, narrow line-width, non LTE distribution of hyperfine intensities, linear and circular polarization, and rapid variability. OH and H20 emission originates from extremely localized regions, often as small as a few A.U. in diameter, in or near some HII regions and some cool, late type stars. The behavior of OH and H20 emission lines is best explained by assuming that the microwave radiation is amplified by some narrow band non-linear process in the sources.

Radio astronomers detected in 1963 hydroxyl molecules (first in absorption), in form of microwave lines at the frequencies 1612, 1665, 1667, and 1720 MHz (1612.231 + 1720.533 (=) 1665.401 + 1667.358), corresponding to a wavelength near 18 cm. Since 1968, some infrared stars have been known to be observable with radio telescopes as sources of such microwave lines, in emission with line widths of 2 kHz or less. In addition to the ground state, a Π state with Λ = 1, the OH molecule has three known electronic excited states, Σ states with Λ = 0, at about 30 000, 60 000, and 90 000 cm−1 above the ground state. These electric-dipole transitions are about 104 times stronger than the magnetic-dipole transitions which produce hydrogen or deuterium lines; the energy of the 2II3/2 state is less than that of the 2II1/2 state what indicates that the observed OH radio emission is subject to far infrared stellar radiation.

Photoelectric photometry of R CrB was obtained during the years 1966–71. During this time the star emerged from a prolonged minimum, showing a change in colour indices which was very like those of the supergiant sequence between G 8 and F 8 if allowance is made for an ultraviolet excess and some interstellar reddening. Superimposed on this rise towards maximum was a sharp minimum, during which the V magnitude was observed to decline through more than 3 magnitudes without any corresponding reddening in the colour indices. Throughout the emergence from the extended minimum the star showed irregular fluctuations of up to 0.2 mag. on a timescale of weeks.

If we believe in the spectral analysis of the atmospheres of RCrB stars these objects should have a pure helium envelope slightly enriched in carbon (SEARLE 1961, DANZIGER 1965). However, they cannot be homogeneous helium stars since in that case they should be on the helium main sequence in the Hertzsprung-Russell-diagram, whereas they are found in the region of late supergiants. Their internal chemical structure must therefore be more complicated than that of a homogeneous helium star. From the viewpoint of stellar evolution the next degree of complexity would be to have a carbon-oxygen core.

It is certain, that the mechanism causing variations of the magnetic field and spectral lines in Ap stars must also cause variations in their luminosities. The light curves are synchronous with the magnetic variations and usually the maximum of the positive magnetic field strength coincides with the minimum of the light curve. In the past the oblique rotator theory was not able to explain easily such brightness change. There is no simple reason to suppose, that the brightness of the surface of a star would increase or decrease at one magnetic pole only. Since that time a few stars were found with some indications for secondary minima and maxima in the light curves, but the first established double wave in a light curve was recently found by H. M. MAITZEN and K. D. RAKOš in HD 125 248 (1970), see Figure 1. It is a very exciting result, only the light curve in yellow light shows two maxima and two minima. The light curves in blue and ultraviolet are very smooth and show no evidence for secondary waves.

Photovisual and spectroscopic observations of HBV 475 in the range λλ 3800–11 000 Å by means of S-20 and S-1 Carnegie Intensifies and prismatic cameras (dispersion 60 A/mm at Hγ), made in 1969 through 1971, have revealed the simultaneous presence of hot and cold sources of radiation, both variable.

In 1969 when the magnitudes were B = 12.45, V = 12.05, U = 11.40 the hot source gave origin to a quite strong ultraviolet continuum and to emission lines among which the most intense are due to: HI (H16 to Hα and P14 to P7); OI 8446; Fell (mult. 27, 37, 38, 42, 48, 49); [NII] 5755; Hel (4471, 5876, 6678, 7065, 10 830 very strong); [OIII] (5007, 4959, 4363); [NeIII] (3869, 3968); Hell 4686 weak.

The peculiar object MHα 328-116 was included in the spectroscopic programme of Asiago since 1965. Its spectral evolution during 1965—66 and 1966—67 has been described in two previous communications (hereafter referred to as Paper I and Paper II) by MAMMANO and ROSINO (1966, 1968).

Identification of approximately 230 emission lines between 3120 Å and 5045 Å have been made on twelve plates taken between 1968 and 1971. The line intensities and ratios fail to give a consistent detailed picture of the nebula, but indicate that on average the electron pressure and density have changed little. The Balmer decrement has become less steep since 1968, indicating higher optical depth. However, the permitted iron lines have completely gone, indicating lower density. Excitation is increasing in some parts of the nebula, as indicated by the increased strength of the [Fe V] lines, and the recent appearance of strong lines of [Ne Y]. The continuum remains weak, though it is easily seen on plates taken on April 8, 1969, and on those taken in June 1971. The radial velocity of the object fluctuates between —50 and — 70 km/sec, and the line widths remain about 120 km/sec. The wing to the violet of the emission lines is still present. The continued brightness of this object, and the relative lack of change in the spectrum over several years, indicate even more strongly that V 1016 Cygni is likely to be a planetary nebula in the stages of formation.

A light variability study in NGC 2264 shows that about 25 % of the pre-main sequence A and F stars show irregular short-period variability and/or brightness changes since 1953. There exists a good correlation between this variability and (other) shell indicators. The scatter in the pre-main sequence band in the color-magnitude diagram can be partially explained by shells and an age spread between 1 and 3 × 106 years is indicated. W 90, a star below the main sequence, has brightened by half a magnitude to V = 12.5 since 1953. The corresponding color variation suggests a ratio of total to selective extinction of about 9.

A photographic survey of variable stars in the Magellanic Clouds is currently in progress at Dunsink Observatory, using photographic material obtained with the ADH telescope. In the course of this investigation a number of Long Period Variables have been measured, most of which lie in the LMC. Several are previously known Harvard Variables (prefixed HV) but some are new discoveries (prefixed DV). The co-ordinates of the new Long Period Variables are listed in Table I.