During the past year we have investigated 72 regions in the galactic plane at 408 and 5000 MHz using the Molonglo and Parkes radio telescopes, respectively.

The angular diameter of a star, combined with observed stellar fluxes in absolute units, allows the absolute fluxes emitted at the surface of the star to be calculated and thus provides a stringent test of model atmosphere predictions. Furthermore, if the flux distribution is observed at all wavelengths contributing significantly to the net flux, an empirical effective temperature can be found. Sufficient data are now available for Sirius (αCMa) to be studied in this way.

Large outbursts from the Sun which are widely extended over centimetre and decimetre ranges and show smooth time variations in intensity have been called microwave type IV bursts (μ-IV). Those bursts generally consist of more than two successive components, the spectra of which move progressively towards the lower frequencies. Furthermore there is a tendency for the decay times, which are defined as the times occupied by the decreasing parts of the bursts, to become successively longer. It seems worth while to make a quantitative study of the time variation of the decay times and flux densities of μ-IV burst spectra and to infer how the physical conditions in the source change with time.

For a number of years the occurrence of isolated groups of apparently related type III bursts has been recognized as a common feature of the dynamic spectral records of solar radio emission at metre wavelengths. Interferometer observations supported the simple hypothesis that bursts of a group were usually located in the same position. In this note we report 80 MHz observations made with the Culgoora radioheliograph which show that although the sources of the bursts in a group tend to overlap one another, their centroids may show marked scatter and that their distribution tends to be along narrow lanes across the solar disk. Since the 80 MHz sources occur at heights (≳ 0.6 R⊙) near or above the plasma level, these lanes doubtless reflect some structural feature of the outer corona.

The theory of stellar convection zones has met three difficult problems. They are non-linearity, non-locality, and non-greyness of radiative transport. Non-linearity is inherent in turbulence. Non-locality is due to the nature of convection, that larger eddies are apt to be more unstable as indicated by the fourth-power dependence of the Rayleigh number on the vertical dimension. Inhomo-geneity and anisotropy of convection, penetration into stable layers, are the consequences. Non-grey radiative transfer through an eddy element is important in a convective atmosphere. In a preceding paper, the author has developed a theory which includes all these effects, using the representative eddy approximation. The emphasis in that paper was put rather in the non-locality and in the radiative non-greyness ; the treatment of non-linearity has been considered mainly in much earlier investigations. In the present note, we shall study how the representative eddy approximation is reconciled with the spectral theory of turbulence, and we shall determine the effective Reynolds and Péclet numbers, the coupling constants appearing in the eddy transport coefficients.

Since its discovery in 1964, interplanetary scintillation has become recognized as a valuable method for investigating the solar wind and the small-scale structure of the interplanetary medium. A particular advantage of the method lies in the ability to study those regions of the medium outside the plane of the ecliptic. To date little has been written about the relation between interplanetary scintillation and solar activity, although regular observations of the source 3C48 during 1965-6 have indicated that a small correlation may exist between the scintillation index and sunspot number. It also appears that anomalous increases in the scintillation index are, on occasion, related to strong flare activity on the Sun.

The 111–110 transition (frequency 4829.65 MHz) of formaldehyde (CH2O) has been recently observed in absorption against a number of continuum sources by Snyder et al. with the Green Bank 140-ft telescope.

Solar X-ray bursts have been observed by a xenon-methane filled proportional counter on the satellite IMP-F. The counter is one inch deep and has a 14 mg/cm2 beryllium window. A slat collimator restricts the field of view to a full width of 5° in the direction of the satellite spin while allowing the counter to see ±60° in the plane containing the spin axis. The effective window area of the counter is 2.2 cm2, after allowing for the collimator.

The U-burst, first identified by Maxwell and Swarup and Haddock, is a type of solar radio event lasting ~ 10 s in which the frequency of the emission at first drifts rapidly downwards, then increases again. On the dynamic spectrum record the burst has the appearance of an inverted letter U.

On 1968 October 30 a large solar flare (importance 3B to 4B) began at approximately 23h42m U.T. We present here a preliminary account of the associated radio emission.

We have used rocket-borne photometers to determine the ultra-violet flux of the full moon in the wavelength range 2400-2900A where atmospheric ozone prevents ground-based observations. The lunar fluxes determined in two rocket firings from Woomera on 9th December, 1965, and 30th August, 1966, are summarized in the Table. The results for the two firings have been corrected to full moon values on the assumption that the lunar phase law is the same at these ultra-violet wavelengths as it is in the visible.

Recently the possibility has been raised of using general relativistic star clusters as models for quasi-stellar sources. The theory of static, spherically symmetric, collisionless star clusters has been developed within the framework of general relativity. In particular, analogues have been found of the Newtonian polytropic models and of Woolley’s truncated Maxwellian systems. However, in view of the importance of rotation on stability in relativistic astrophysical problems, it is of considerable interest to include the effect of rotation in relativistic stellar dynamics.

As shown recently by Y. Osaki super-massive stars with mass M < 3.5 × 105M⊙ can, in the absence of rotation, reach the hydrogen-burning main sequence before the onset of general relativistic instability. Such objects are then pulsationally unstable. A considerable simplification is introduced if one considers only very massive stars, for which the relative amplitude of the fundamental mode of oscillation is practically constant. This sets a lower limit of 104M⊙ to the mass that can be considered. The upper limit is also reduced to 2 × 105M⊙ if one neglects the relativistic correction. One necessary step in the study of non-linear oscillations of massive stars is to derive a differential equation for the adiabatic pulsations. The relativistic correction could be taken into account in the following way.

The heights of solar radio sources at 1424 MHz and 696 MHz have been measured during the years 1965 and 1966. Solar activity at this time was near minimum. The number of radio sources appearing on the solar disk rarely exceeded three at any time and it thus was possible to resolve the majority of these with a high resolution grating interferometer. Many of the previous height measurements at these frequencies have been made near times of maximum solar activity and the confusion of sources within the beam has limited their accuracy. The number of sources studied here is quite considerably higher than in any previous investigation at these frequencies, and the period of observation has been continuous.

Regions of the interplanetary medium currently inaccessible to space vehicles may conveniently be studied using the radio scattering properties of the interplanetary plasma. These effects may give rise to angular broadening of radio sources sufficiently close to the Sun, or to amplitude scintillation of sources of small angular size.

The thermal radio-frequency emission, Ev, per unit volume per unit frequency interval may be computed (e.g. from equations given by Oster) and compared with the volume emissivity E(Hβ) in the optical Hβ transition.

Dodson and Hedeman discovered an unexpected effect in the occurrence of solar proton events as revealed by polarcap absorption (PCA). When the 48 events in Bailey’s Catalog of the Principal PCA Events, 1952-1963 are distributed with the phase of the moon there is a gap of several days near full moon; also, many more events occur when the moon waxes than when it wanes. Dodson and Hedeman did not find similar, apparent departures from random distribution either with a mean solar rotation period of 27.3 days or for solar flare events. They concluded that ‘at the present time it is not clear whether the 29.5 day “effect” is related to the sun or the moon or is only a statistical accident’.

Recent high resolution Ha photographs by Bray and Loughhead (unpublished) have confirmed that much of the chromosphere is quite non-uniform. Rather, it appears to consist of many distinct structures, which have a typical transverse section of several seconds of arc, i.e. about 2000 km.

The aim of the research reported in this contribution is to attempt to analyse the non-uniform chromosphere by first investigating a single structure. A simple model of a structure was chosen. From the model, line profiles of Hα and the H and K lines of Can were computed assuming a variety of physical conditions to exist in the feature.

A natural introduction to this topic is a brief discussion of two phenomena observable in interplanetary space near Earth—the energetic storm particle event (ESPE), and the recurrent Forbush decrease.

In an IAU Symposium report, Spiegel has made mention of the basic reasons for the astrophysical interest in the problem of non-linear convection of a horizontal fluid layer heated from below. Unfortunately the two important parameters of this problem—the Prandtl number P and the Rayleigh number R take on ‘extreme’ values in the astrophysical situation and this presents immediate difficulties as far as the numerical solution of the basic equations go.