Radiation pressure due to absorption in spectral lines can play an important role in the equilibrium conditions in the outer portions of the stellar atmosphere. In particular, it is suggested that spicules ejected from the upper solar chromosphere as well as other fine structural features of the chromosphere may very likely be a result of radiation pressure from Ly-a. Numerical computation of the radial radiation pressure due to Ly-α in the chromospheric model of Vernazza et al. (1973) gives a maximum outward pressure gradient due to radiation that is 0.06 of the gravitational pressure gradient. The model chromosphere of Vernazza et al. is a spherically symmetric model that reproduces the average Sun Ly-α intensity and profile. Supergranule borders (network), where spicules arise, are known to have a Ly-α brightness some ten times the average Sun brightness. It is suggested, therefore, that in the network areas the pressure gradient due to radiation is comparable to (and oppositely directed to) that due to gravitation and that in local areas of unusual brightness the net pressure gradient is outwards. An estimate of the expected outwards velocity of spicule material driven by radiation pressure was made by equating the kinetic energy flux (½JϱV3) to the Ly-α energy flux. This gives values of v that are comparable to observed spicule velocities. Also, the maximum outwards radiation pressure occurs in the upper chromosphere where the material density in the model of Vernazza et al. is close to the value 4 × 10−14 g cm−3, which is about a factor of two less than the currently accepted value for spicules. Increased density in the network areas could easily remove this minor difference.

The solar limb has been studied visually and photographically in the light of Hα and the combined results are reported. The appearance of a spicule changes with wavelength settings. Effects due to changes of spicule velocities with height and opacity-wavelength relations are indicated. A number of bright spicules seen in the wings are identified with bright mottles seen in the line core. The dark band appears between Δλ=± 1.0 Å and Δλ= ±0.70 Å. Examination of a small mottle cluster in a series of photographs taken in steps from Hα+ 0.85 Å to Hα+ 0.10 Å leads to a three-dimensional model of a typical mottle cluster.

The subject of this symposium is the fine structure of the solar chromosphere. Progress in this field of reserach will depend to a high degree of the quality of seeing, resp. on the effective angular resolution available on the ground. Today's situation of solar ground seeing has changed distinctly in the last years. I would like to report here a few new aspects, which could be condensed into 3 questions:

1. (1) Are there on the ground ‘good seeing windows’, comparable in quality with stratosphere results obtained from balloon borne equipment?

2. (2) Is there a chance to resolve from the ground the solar scale height, corresponding to about 0.1″.

3. (3) Is there a residual fundamental atmospheric seeing noise resp. a basic limit to the atmospheric Modulation Transfer Function (MTF)?

We are reducing a set of spectra covering the region from 3400 to 4330 Å that show both spatial and spectral structure in chromospheric emission lines from many elements and ions. The spectra were taken with the vacuum tower telescope at the Sacramento Peak Observatory with the spectrograph slit tangent to and touching the limb. Thus some height resolution in the chromosphere is present. Many of the lines show doubly-reversed emission, often asymmetric either for the whole line or for fine structure. The quality of the data is such as to improve our understanding of line blends. The spectral structure can be compared with the well-known structure in the H and K lines of Ca II and in the resonance lines of MgII in order to deduce a model for the lower chromosphere.

The fine structure in the chromosphere is probably controlled by magnetic flux concentrations which cellular convection produces in the photosphere. We will first deal with this flux concentration by the convection and its constraints. Then we will discuss the chromospheric network and its main constituent, the spicules. Following that we will look at the inside of the network, i.e. the fibrils.

A number of observers have noted the presence of bright structures near the cores of the chromospheric rosettes when observed in the far wings of the Hα line (eg Hα ±7/8 Å). Dunn and Zirker observed these bright structures with the highest possible resolution using the Sacramento Peak vacuum solar telescope. They find that these bright regions exhibit a very intricate fine structure which can be followed out much further into the Hα line wing (eg Hα + 2 Å) and even into the continuum. They called this fine structure ‘solar filigree’, the name referring mainly to the collective appearance of the fine structure elements. The elements themselves appear as dot-like structures and frequently also as small wiggly structures called ‘crinkles’. The properties of the filigree structure are summarized as follows:

1. (i) Size: Measured diameter of the crinkles and dots equals 0.25, 0.40 and 0.60″ at Hα + 2 Å, Hα ± 7/8 Å and Hα ±5/8 Å respectively. The telescope resolution equals 0.22″ so that at Hα + 2 Å the structure is extremely small. The drawings in Figure 1 show typical sizes of the crinkles and network patterns in the filigree.

2. (ii) Contrast: Filigree is enhanced in the blue wing of the Hα line. Measured contrast, uncorrected for seeing, equals 5–10%.

3. (iii) Relation to the Granulation: The filigree structures tend to lie between the granules. This is, however, not a strict rule. It seems that in the course of their lifetime the granules move the filigree structures around with velocities of about 1.5 km s-1. Some of the crinkles also seem to wash out temporarily until compressed again by a new granule. The detailed structure of the filigree, therefore, changes significantly over times comparable to the granule lifetime. The overall structure is, however, preserved over much longer periods of time. The granulation pattern when observed in the continuum well outside the Hα line appears very peculiar in that it has substantially decreased in contrast. It appears ‘soft’ similar to granulation washed-out by seeing. This abnormal granulation can be traced over long times (> 30 min) and coincides in location to the filigree location. It is, therefore, definitely real.

4. (iv) Relation to the spicules: The filigree structure falls near the center of the Hα chromospheric rosettes. These rosettes consist of dark elongated mottles which should probably be identified with spicules. There is, therefore, at least a coarse relation between the occurrence of spicules and the filigree. There is no clear evidence that variations in the filigree pattern are related to the generation of spicules. Some spicules seem to originate from the spaces between the crinkles. Too few, however, to conclude a definite relation.

5. (v) Relation to the magnetic field: Beckers studied the filigree with the Universal Birefringent Filter in the magnesium b1 and b2 lines. It is very well visible in the far wing of the lines (eg. b1 ±0.8 Å). When traced into the line core the structures increase somewhat in size, as they do in Hα, and form structures similar to, and perhaps identical with, the so-called photospheric network. In the magnetically sensitive b2 line one sees a one-to-one correspondence between these network structures and the magnetic field so that, at least in the layers seen near the core of the b2 line, there is a one-to-one correspondence between the filigree structures and the enhancements in the magnetic field. Simon and Zirker (Solar Physics, submitted for publication) using a spectrograph also found that the filigree occurs in regions of enhanced magnetic field. However, in contrast to the filter observations, they found the magnetic field regions to be much more diffuse (2–3″) so that there is not a one-to-one spatial correspondence between filigree and magnetic field structure.

I have tried to understand whether spicules appear dark or bright against the disk in Hα by comparing high resolution pictures in Hα and Ca II K. Since the chromospheric structure is more apparent near active regions, I have first compared the Hα and K fibrils near an active sunspot; these are illustrated in my review paper (this Symposium, p. 161). The elongated fibrils in the active center appear similar to spicules, but much larger and more horizontal. K line pictures show bright features better, while Hα shows dark features better.

We describe time-series observations of small-scale Can emission features located outside the network in the quiet chromosphere. Simultaneous spectra in K and λ8542 show unambiguously that the evolutionary behaviour of the K-line profile is due to an outwardly propagating velocity pulse. Assuming that this pulse is a progressive acoustic wave, as suggested by the inferred flow parameters, we show that the wave loses mechanical energy in traversing the chromosphere. This implies that the bright Can features (K grains) are the manifestation of local heating in the chromosphere, possibly by shock waves.

The line formations of Hα and CaII H and K are compared in order to differentiate the various mechanisms giving rise to observable contrasts in the emergent intensities. Table II summarizes the criteria for distinguishing between horizontal spatial variations in temperature, density, and turbulent velocity.

The Hat Creek Observatory's two-element interferometer and the NRAO 3–element interferometer have been used at 1.3, 3.7– and 11.1–cm wavelengths respectively, to study the fine structure of the radio emissive regions on the Sun. Observations of the quiet Sun at 1.3 cm show sudden increases followed by a gradual decrease in the fringe amplitude lasting for typically about 5–8 min. Assuming these events are identical in nature, a plot of peak amplitude against the projected baseline at the time of the event suggests emission from a region of angular size of about 10″. The corresponding brightness temperature is 50000 K. It is possible that these events may be related to the appearance and disappearance of groups of spicules or mottles.

The paper contains a review of attempts to understand observational data on the inhomogeneous structure of the upper solar chromosphere, particularly the spicules. Observations over the electromagnetic spectrum from millimetric to the ultraviolet give equivocal results; and in spite of recent progress, models for these inhomogeneous structures still face difficulties in accounting for the observed radiation. New observations in the visible are needed in order to answer even the simplest questions of kinematics, while the interpretation of the line spectra still poses many difficulties. The acquisition of EUV data at high spatial resolution should lead to major progress in the field, just as it already has done in showing the EUV emission to be concentrated over the network boundaries.

Further observations of limb and disk intensity ratios of emission lines in the EUV solar spectrum were obtained on a Skylark rocket flight on 5 August 1971. Analysis of the data has shown that the observations support the existence of a steep rise in temperature in the transition region between Te ~ 6 × 104 K and 3 × 105 K. The average absolute height of the transition region above the visible limb has been measured with a greater accuracy than previously possible and is 1700 km ± 700 km. An independent method using arc-length measurements of spectrum lines gives an absolute height of 2100 km ± 850 km. Emission from lines which are optically thick in spicules is observed to extend to heights of 10000 km above the transition region. The observed decrease of Fen emission with height is consistent with current spicule statistics.

The emission measure Ne2 dh, which is so useful in reducing XUV as well as radio data, is redefined as f(T) dT where f(T) is called ‘thermal emission measure’. Theoretical predictions for f(T) on the basis of a simple, one dimensional, steadily expanding atmosphere are presented. Depending upon the boundary conditions, and essentially upon the mass flux, two very different behaviours show up.

1. (1) With a mass flux compatible with an extrapolation of the solar wind flux, f(T) would correspond to a transition region controlled by a constant conductive flux. It can be fitted, as was shown e.g. by Athay (1966), with the XUV observations of the integrated disk for lines formed at temperatures higher than about 2 × 105 K.

2. (2) But only with a mass flux enhanced by a factor of 50 to 100, can we interpret the radio spectrum and the XUV observations of lines formed at temperatures less than 2 × 105 K.

The suggestion is made that a single model can reconcile both behaviours: below 2 × 105 K the emitting region is squeezed so that it covers 1–2% of the surface that it occupies at higher temperatures where all structures begin to merge, filling the whole corona. This variation with height compares with usual models derived from spatially resolved observations; as usual the underlying magnetic field is expected to support this channelling.

The basic features of the so-called ‘cloud-model’ (Grossmann-Doerth and von Uexküll, 1971) are discussed. The model permits us to derive approximate values of certain properties of the elements of the chromospheric fine structure by comparing the observed profiles of strong Fraunhofer lines with a simple theoretical formula involving four free parameters. The method is only applicable if two basic conditions are met: (1) The absorption coefficient varies with wavelength according to a gaussian with a broadening parameter that is independent of height, (2) the variation of the source function Sv with height is sufficiently small so that may be replaced by a constant. The validity of these assumptions, in particular the latter one, is discussed. It is shown that in the case of Hα there are reasons to believe that the assumption is valid provided the features are not very close to the limb (sin θ = 0.9 may constitute the border line). There is also some evidence to suggest that the cloud model will not work for the CaII K-line.

This review concerns itself with the measurement of the effects of chromospheric motions and the diagnosis of those motions themselves over approximately the last ten years. The different types of observational techniques are described. The different size regimes of motions are reviewed and their possible effects on observable quantities are discussed. The different types of motion in the lower chromosphere are reviewed, including microturbulence, ‘layered’ motions, and observations and interpretation of asymmetries in the core of the CaII K line. The observation of motions in Hα mottles on the disk and spicules on the limb are reviewed, from the standpoint of both line profile analysis and broad-band measurements. The interpretation of these motions and attempts to construct empirical models of chromospheric fine structures are discussed.

Oscillations and waves have been observed in Hα in sunspot umbras and penumbras. They appear to be basically of the Alfvén type, but do not carry away sufficient energy to account for the sunspot energy deficit.

There is a pronounced difference in the velocity structures of supergranule centres and boundaries. In the former it is oscillatory, of shorter period but bigger amplitude than in the photosphere. In the network there are oscillations of the fibrils but, on the whole, little evidence of wave motion along their lengths except perhaps near their diffuse ends.

In plages, there are velocity oscillations which at times may be phase coherent over extensive regions up to 105 km apart; there are suggestions of a connection between these times and the occurrence of flares. There is also a report of a velocity pulse travelling along an active-region fibril with a velocity of propagation of 150 km s-1.