The 5-min oscillations in Hα plages as reported by Bhatnagar and Tanaka (1971, Solar Phys.24, 87) are discussed with respect to spatial – both horizontal and vertical – phase relation. Data for analyses include (a) Hα filtergrams at −0.5 Å (1050 min in total) (b) Hα filtergrams simultaneously taken at ±0.5 Å (60 min) (c) Simultaneous Hα−0.5 Å and FeI 5233 Å videomagnetograms (Doppler mode) (100 min in total). Projection (16 and 35 mm) -photocell method (Bhatnagar and Tanaka, 1971) was used for tracing the data with special attention to the guiding and reproducibility. Also photographic prints were used for comparison. The results are as follows:

1. (1) those showing oscillatory brightness changes are dark dot-like features (size ≲ 1″) in brighter regions of plage and slightly elongated features (1 × 1.5–2″) in less bright regions. The mean number density of these features is one per 2200 ± 500 km square.

2. (2) There are always a few regions with a size 8″ or more in one plage which show very regular oscillations with a period of about 5 min for a long time (at least 6.5 h, Figure 1). The distribution of the intervals between successive intensity peaks is much narrower than the Gaussian distribution and has a skewness towards longer period. The other places in the plage show relatively irregular patterns, but with the mean interval of the fluctuations still being 5 min.

3. (3) In regular 5-min brightness oscillations the phase of oscillations in one wing is complementary to that in the opposite wing, indicating velocity oscillations. When the intensity records are relatively irregular there are many cases (76%) in which the blue wing patterns can be matched well for the whole time sequence with the red wing records by advancing the red wing records by 196 ± 18 s, (see Figure 2) indicating that the observable dark feature first appears in rising mode and then with a constant time lag (close to the free fall time from −0.5 Å to +0.5 Å) appears in the falling mode.

4. (4) The records of brightness variations are very similar when the aperture is changed from 4″ (minimum reliable aperture) to 10″ in most places. Comparison of the oscillation phases in all the measured points in one plage shows that a plage can be divided into several areas (for a certain time) with sizes from 7000 km to 30000 km in each of which the phases are very close to each other for 2 to 10 cycles. Figure 3 shows the distribution of the intensity peaks of oscillations measured at Hα–0.5 Å (the abscissa is time, the ordinate indicates space points arranged geometrically). The boundaries of these areas vary after the coherent cycles. The iso-intensity curves of the intensity-time records (see Figure 3) have slopes corresponding to horizontal phase velocity 100 km-600 km s-1.

5. (5) The phase relation between Hα regular 5 min and FeI 5233 Å 5-min velocity oscillations at the same place is almost constant for more than 5 cycles although the phase lag is different from place to place. (See, for example, the period from 20:00 UT to 21:00 UT in Figure 1. The smoothed-out curve denotes the record of the velocity oscillations in λ 5233 Å taken simultaneously.) The distributin of the phase lags is peaked at about 1 min. Assuming 1000 km for height difference of two lines the vertical phase velocity is equal to 17 km s-1.

6. (6) There is an indication that the flare originating in the plage occurs about 20 min after the oscillations start to be in phase (1-min accuracy) in very large scale (≳20000 km) although the samples (6 flares in three different plages) are not large enough for statistically meaningful discussions.

In the presence of a magnetic field, three types of magnetoatmospheric waves - magnetoacoustic mode, magnetogravitational mode, and hydromagnetic mode - can propagate in a stratified atmosphere, in contrast to the propagation of two types of atmospheric waves - acoustic mode and gravitational mode - in the absence of a magnetic field. The exact manner of propagation of the magnetoatmospheric wave is extremely complex, and most studies have been confined to certain specific circumstances, such as an isothermal atmosphere permeated by a uniform magnetic field (McLellan and Winterberg, 1968; Bel and Mein, 1971; Michalitsanos, 1973), and atmosphere in magnetohydrostatic equilibrium (Yu, 1965; Chen and Lykoudis, 1972) and the propagation across a density discontinuity (Stein, 1971).

Every sunspot group is surrounded by regions of weaker magnetic field; traditionally those parts which appear in emission in resonance lines are called plages. There is a more extensive region of enhanced magnetic field which is not bright because the magnetic field is horizontal; this region is easily discernible by dark horizontal fibrils in Hα and is included in our discussion. The plage includes most of the magnetic field in an active center and is the locus of most of the flares.

Weakenings of 14 FeI lines observed in photospheric faculae by Chapman and Sheeley (1968) have been analyzed using Unno's (1956) LTE theory of Zeeman triplet formation. These observations can be explained by a facula-photopshere temperature excess of ≲250 K in the layers ~ 100-400 km above τ5000 = 1. The inferred magnetic field in the faculae is ~ 1000 G.

In a recent paper Stellmacher and Wiehr (1973) discussed in detail the discrepancies between ‘continuum-’ and ‘line-profile facula models’. Their calculations show that continuum models (e.g. Kuz'minykh, 1962; Schmahl, 1967; Chapman 1970; Wilson, 1971) encounter the problem that the facula contrast increases towards the limb but decreases with wavelength λ. This discrepancy evidently cannot be due to NLTE effects since it concerns the continumm data.

Using the high resolution photographs of the photosphere taken at the Ondřejov observatory a comparison of positions and forms of some chromospheric features with the distribution of photospheric fine structure morphological elements in the August 1972 active region has been made. The close relationship between active filaments, chromospheric threads, fibrils and photospheric small nuclei, dark interpenumbral and intergranular spaces or rudimentar penumbra is shown. Some considerations are demonstrated with the aid of computed models of the horizontal component of the magnetic vector simulating the real distribution of magnetic fields in the studied region.

The relatively cool Hα emitting regions of solar prominences are evidently surrounded by an interface or structure of intermediate temperature that separates them from the surrounding hot corona. At present, very little is known about this prominence-corona interface (hereafter called the PC interface), because most of its distinctive radiation is in the EUV where observations of high spatial resolution have not been available. In this paper we report on studies made on photometrically calibrated monochromatic images of prominences and the surrounding corona recorded on slitless spectra made during the eclipse of 1970, March 7. These were obtained with a rocket-borne Wadsworth spectrograph flown by the consortium group (Speer et al., 1970) and cover the range λλ850–2159 with a spectral resolution of 0.17 Å and a spatial resolution better than 10″ over half the spectrum. They include spectral lines of a number of ionic species from Hi through Nixv. Because each species exists in a rather narrow temperature range, some of the monochromatic images show the low temperature prominence itself; some the coronal structure surrounding it; and some the intermediate structure we have called the PC interface.

An example of an erupting filament that exhibited helical structure during a flare on 1971, September 17, was discussed. This sort of event is not rare. Examples of rotational effects in flares, prominences and surges have been reported by others such as Öhman (1968). Accompanied by type II radio emission and the visible flow of Hα gas across 1/2 the Sun's surface, the flare was very interesting in itself. However, here the emphasis is on possible ways in which the observed twist might develop. The latter is important because it yields information about: (1) the state of the subphotospheric field itself, and (2) the possibility of storing flare energy in the twisted filaments.

Observations prompted by Sheeley's discovery of outflowing CN bright points from sunspots have established that there occur small moving magnetic features (MMF's) near sunspots. These MMF's exhibit a highly ordered pattern of movement directly related to the associated sunspot. The observations are consistent with the concept of magnetic flux outflow (MFO), a process whereby net magnetic flux of the same polarity as the sunspot is transferred from a decaying sunspot to the surrounding magnetic network. Small magnetic flux concentrations are apparently convected outward by a velocity cell centered on the sunspot. Doppler spectroheliograms have provided evidence for such systematic outward velocities extending as far as 10000 to 20000 km beyond the outer edge of the penumbra of some sunspots, which is comparable to the extent of MFO. While MFO is best observed by means of time-lapse movies of the magnetic fields, it is also manifested morphologically on individual magnetograms by features that resemble moats (or bays) and wreaths around only those sunspots where MFO is present. Two examples of magnetic features streaming toward and into rapidly forming sunspots are described, providing evidence for the occurrence of magnetic flux inflow (MFI) associated with the growth phase of sunspot development. It is, therefore, likely that MFI and MFO are basic aspects of the evolutionary development of sunspots. Observational and instrumental aspects relevant to the investigation of MMF's are described in the Appendix.

Bumba (1963) inferred from plots of areas of sunspot groups against time that stable spots may show a phase of slowest decay which is independent of the spot's area. From the Greenwich Photoheliographic Results we analysed the areas of individual spots. Indeed there is a clear lower limit in the decay rates close to Bumba's value dA/dt ≃ −4.0 × 10−6 visible hemisphere/day However, only a small fraction of regular spots and an even smaller fraction of the irregular spots pass through the phase of slowest decay.

In this paper we investigate the physical processes that lead to the growth and decay of magnetic flux in and near sunspots.

An initial phase of rapid growth is characterized by the emergence of magnetic flux from the deep convection zone. As the flux rope rises through the surface the magnetic field is swept to the junctions of the supergranular network where sunspots are formed. These flux concentrations follow the footpoints of the emergent flux rope as they rapidly move apart.

Evolution of chromospheric fine structures is discussed from observational points of view. First, various lifetime determinations are reviewed. Then chromospheric oscillation in the interior of supergranule cell is discussed in relation to chromospheric ‘grain’ structure. We show that the grains are collectively associated with the 180-s oscillations as well as that in K line the grain is identified with local disturbance generated at lower level and propagating to upper levels. Finally the details of evolution of individual mottles are presented. We show some cases which indicate the propagation of a disturbance along the mottle axis. Statistics of evolution patterns of mottles implies a significant tendency for a system of double dark mottles encasing a bright mottle.

We present off-band Hα filtergrams showing the bright Hα network (solar filigree) and its evolution, with subarcsecond resolution. In quiet regions the total area covered by elements in the network is constant over periods of hours but changes in the geometry are observed over 10 min (the timescale of granule and dark mottle lifetimes), that can be interpreted as due to (macro)turbulent horizontal velocities of about 1 km s-1. The network can be resolved into strings of tiny bright dots; on the order of 100 per supergranule cell boundary. These dots may mark the location of magnetic filaments. The contrast of the features is higher in the blue than in the red wing which may indicate that in the lower chromosphere, the downdraft at supergranule boundaries does not occur in the bright crinkles themselves but rather adjacent to them.

Series of simultaneous spectra in different spectral regions, representing all ‘visible’ layers of the solar atmosphere, of high spatial and temporal resolution are used to compute power spectra and diagnostic diagrams of intensity and velocity fluctuations in order to study the dynamic response of the chromosphere to disturbances from below.

The problem of the solar atmospheric heat and energy balances is here divided into three sections.

1. (i) Basically the problem is an outward temperature gradient involving additional radiation losses whose estimate involves a model atmosphere. Proposed models described are based on optical and radio emissions together with various theoretical considerations such as thermal conductivity. The models are useful in providing estimates of the required flux of mechanical energy from below. However, they are limited by the use of averaging processes, and the failure to take full account of the many structural features which make up the chromosphere.

2. (ii) It has long been accepted that the main energy input is by acoustic waves generated in the ‘small-bubble’ Vitense model convection zone. We review recent observations of chromospheric magnetic and velocity fields and their interpretation which strongly suggest that this model is not valid and that the acoustic theory of heating is without basis. At the same time the new data remove the objections of Osterbrock and others to a theory of heating based on Alfvén and/or slow-mode hydromagnetic waves.

3. (iii) Observations of a variety of individual magnetic-plasma chromospheric structures, together with the adoption of heating by waves which follow the field lines, suggest a new approach to the whole problem of the heat balance. We discuss the heating problem in a number of these magnetic-plasma structures including the emission network, spicules and related disk features, arch filament systems and flares.

The relation between local criteria for hydrodynamic instabilities and stellar stability is reconsidered. This leads to a discussion of the classical problem of stellar pulsation with a special emphasis on the pressure modes of high order spherical harmonics. It is shown that within the frame of the geometrical approximation, acoustic rays exist which are trapped inside any specified slab of a stellar envelope. The pressure modes associated with such rays are conjectured to give rise to small scale instabilities whenever a region exists in an envelope where any local criterium indicates pulsational instability. The model is expected to support and somewhat extend the recent developments concerning the role of pressure non-radial modes on the dynamics of the solar atmosphere and the heating of external layers.

Preliminary report available from Observatoire de Nice.

The structure of the quiet solar atmosphere is dominated by effects due to the magnetic field distributions produced by field concentrations at the supergranual boundary regions. This effect has been studied by previous workers (Kopp and Kuperus, 1968, Kopp, 1972), and is also evident in the enhanced intensity at the SBR's of spectral lines formed in the upper chromosphere and transition region. An attempt is being made to evaluate the field distributions more precisely, taking into account the physical processes involved in the interactions between the field and plasma, and thereby to predict the intensity distributions expected for the EUV spectra. This work is not yet complete, but some important results can be foreseen.

Excess heating of the active solar atmosphere is interpreted as the decay of MHD slow-mode waves produced in the corona through the non-linear coupling of Alfvén waves supplied from subphotospheric layers. It is stressed that the Alfvén-mode waves may be very efficiently generated directly in the convection layer under the photosphere in magnetic regions, and that such magnetic regions, at the same time, provide the ‘transparent windows’ for Alfvén waves in regard to the Joule and frictional dissipations in the photospheric and subphotospheric layers. Though the Alfvén waves suffer considerable reflection in the chromosphere and in the transition layer, a certain fraction is propagated out to the corona. A large velocity amplitude, exceeding the local Alfvén velocity, is attained during the propagation along the magnetic tubes of force into a region of lower density and weaker magnetic field. The otherwise divergence-free velocity field in Alfvén waves, in such a case, couples with a compressional component (slow-mode waves) which again is of considerable velocity amplitude relative to the local acoustic velocity when estimated by using the formula for non-linear coupling between MHD wave modes derived by Kaburaki and Uchida (1971). Therefore, the compressional waves thus produced through the non-linear coupling of Alfvén waves are eventually thermalized to provide a heat source. The introduction of this non-linear coupling process and the subsequent thermalization of the slow-mode waves may provide a means of converting the otherwise dissipation-free Alfvén mode energy into heat in the corona. The liberated heat will readily be redistributed by conduction along the magnetic lines of force, and thus the loop-like structure of the coronal condensations (or probably also the thread-like feature in the general corona) is explained in a natural fashion. Matter density in the arches is expected to be higher than the ambient density because of the higher temperature and, consequently, greater scale height in the arches.

The stability of a magnetic flux rope is investigated in the presence of an Alfvén wave flux. It is shown that instabilities will arise when the transverse perturbation velocity exceeds the Alfvén velocity for the flux rope. With this restriction the maximum permitted Alfvén energy fluxes have been calculated under a variety of conditions typical of faculae, flares and sunspots. It is shown that the power requirements for faculae are very close to the theoretical limit and that flares may or may not exhibit a flash phase depending on whether the stability criterion is exceeded or not. The Alfvén fluxes which have been suggested for sunspots are, however, generally less than the limit based on this criterion.