On the sunspot structure

The fine structure of a sunspot is studied on a series of photographs obtained during the third flight of the Soviet Stratospheric Solar Station. The main results are as follows:

1. The micro-photometer tracings on the frames show extremely high Rayleigh resolution of small elements, the smallest distances being near to the theoretical limit. The half-widths of the brighter elements are given in Tables III and VI. The corrected brightness of umbral dots has large dispersion.
2. The dimensions of the smallest dots are equal to the diffraction image of bright points. So the real radii of these objects are smaller than 150km, which is consistent with opaque models of sunspot umbra.
3. The penumbra and umbra structure (dark and bright objects) is in good agreement with the picture of magnetic field splitting in a system of magnetic ropes giving rise to the magnetic arcs in the chromosphere and corona. Only in the umbra do we meet the large scale continuities.

     

Étude morphologique et cinématique des structures fines d'une tache solaire

A sequence of 34 photographs of the main spot of the group H 26 (Daily Maps of the Sun, Freiburg 1970, Rome number 5847) has been obtained with the 38 cm refractor of the Pic-du-Midi Observatory, showing throughout a resolution very close or equal to 0′'.3. An interval of 3 hr is covered. The pictures taken at intervals of 6 min approximately permit to study the fine structure of the penumbra and associated phenomena:

1. The penumbra appears to consist of bright grains, generally lined up in the form of filaments, showing up against a dark background (see Figure 1).
2. The bright grains form all over the penumbra (see Figure 5).
3. They move toward the umbra of the spot. Their horizontal velocity is zero at the border penumbra-photosphere and maximum at the umbral border (0.5 km s^?1) (see Figures 3,4 and 8). Therefore, the grains never originate in the photosphere nor do they enter it.
4. They disappear in the penumbra proper or, if they form near enough to the umbra and live long enough, they can enter the umbra and their appearance becomes similar to that of umbral dots.
5. The life time of the grains is a function of their place of origin within the penumbra: It is maximum and of the order of 3 hr or more for those forming in the middle part of the penumbra, and 50 and 40 min respectively for the points formed in the inner and outer part of the penumbra.

     

X-Ray and microwave observations of active regions

Radio and X-ray observations are presented for three flares which show significant activity for several minutes prior to the main impulsive increase in the hard X-ray flux. The activity in this ‘pre-flash’ phase is investigated using 3.5 to 461 keV X-ray data from the Solar Maximum Mission, 100 to 1000 MHz radio data from Zürich, and 169 MHz radio-heliograph data from Nançay. The major results of this study are as follows:

1. Decimetric pulsations, interpreted as plasma emission at densities of 10⁹–10¹⁰ cm^?3, and soft X-rays are observed before any Hα or hard X-ray increase.
2. Some of the metric type III radio bursts appear close in time to hard X-ray peaks but delayed between 0.5 and 1.5 s, with the shorter delays for the bursts with the higher starting frequencies.
3. The starting frequencies of these type III bursts appear to correlate with the electron temperatures derived from isothermal fits to the hard X-ray spectra. Such a correlation is expected if the particles are released at a constant altitude with an evolving electron distribution. In addition to this effect we find evidence for a downward motion of the acceleration site at the onset of the flash phase.
4. In some cases the earlier type III bursts occurred at a different location, far from the main position during the flash phase.
5. The flash phase is characterized by higher hard X-ray temperatures, more rapid increase in X-ray flux, and higher starting frequency of the coincident type III bursts.

     

Space and time variations of the solar Na D line profiles

Preliminary results are presented of observations of the solar Na D lines obtained with high space and time resolution (2.4″ × 2.4″), (6 s). The following conclusions may be drawn.

1. The line profiles vary strongly with space and time implying that time averaging over a long period and large area will not produce the ‘true’ profile.
2. The centre-limb increase in apparent Doppler width in the D lines is intrinsic. It is not due to space or time averaging.
3. The amplitude of the 300-s oscillation may range up to 1.5 km/s in the region of formation of the D lines. Large line asymmetries are associated with this motion. Observations which do not resolve this motion can not be considered adequate.
4. The variation of the D line profile caused by the 300-s oscillation may be described as follows: (a) The core is raised and lowered without change of shape, (b) The wings broaden as the central intensity rises and narrow as it falls. These variations are qualitatively explained by the scanning of the line formation region through the solar atmosphere.
5. Doppler width values derived from pairs of D line profiles are strongly correlated with the motion of the element observed. Hotter elements move upward, cooler downward.
6. Indications of running waves have been found in the time variation of the core line bisectors.

The profile variations observed provide a framework in which various properties of the centre limb variation of these lines may be considered. In particular they show that any expectation of accuracy in profile coincidence above a certain value must be doomed by the intrinsic variability of the solar atmosphere.     

Spectral analysis of four quiescent prominences observed at the Peruvian eclipse

Two-dimensional distributions of kinetic temperature, density and turbulent velocity are obtained for four quiescent prominences observed at the Peruvian eclipse of 12 November, 1966.

1. The kinetic temperature derived from line widths is around 6000–7000 K in the central part of prominences and rises to 12000K in both edges and possibly in the top of prominences.
2. The turbulent velocity shows a similar tendency, being 7–9 km/sec in the central part and ≈ 20 km/sec in the outer part. The turbulent velocity also increases slowly towards higher heights in the prominence.
3. The electron density derived both from the Stark effect and the intensity ratio of the continuous spectra turns out to be about 10^10.2–10^10.6 cm^?3 in the central portion of two prominences.
4. From the width and the intensity, neutral helium lines are shown to originate in the same region as hydrogen and metallic lines where the kinetic temperature goes down to 6000 K. This indicates that neutral helium is emitted after the ionization due to UV radiation from the corona and the transition region.

     

The evolution of the photospheric network

A time-lapse sequence of spectroheliograms in the bandhead of CN at λ3883 reveals the following behavior of the photospheric network with time:

1. There is a steady flow of bright ‘points’ (? 1000 km in diameter) laterally outward from sunspots at speeds on the order of 1 km·sec^?1. After traveling about 10 000 km from a sunspot they either conglomerate to form fragments of the photospheric network or disappear.
2. Spatial changes in the network pattern seem to take place by means of the shifting of network fragments laterally on the solar surface. Although most small-scale details are recognizable after 5–10 minutes, within 30 minutes nearly all the details have changed completely. In contrast to this, the large-scale network pattern seems relatively unchanged after 2 1/2 hours.
3. Occasionally ‘new’ network, not resulting from the lateral motion of bright features from either previously existing network or sunspots, appears on the solar surface. This process consists of the formation in approximately 10 minutes of bright points and a darker-than-average feature between them. The dark feature disappears in another 5–10 minutes and the bright points separate at a relative speed of a few km·sec^?1. If the event is of a sufficiently large magnitude, a sunspot will appear.

These observed changes of the photospheric network with time are interpreted as formation and motions of photospheric magnetic fields. It is suggested that these motions reflect the presence of both short-lived small-scale and long-lived large-scale photospheric currents such as one might expect from the granulation and the supergranulation.     

Delta spots and great flares

Using eighteen years of observations at Big Bear, we summarize the development of δ spots and the great flares they produce. We find δ groups to develop in three ways: eruption of a single complex active region formed below the surface, eruption of large satellite spots near (particularly in front of) a large older spot, or collision of spots of opposite polarity from different dipoles. Our sample of twenty-one δ spots shows that once they lock together, they never separate, although rarely an umbra is ejected. The δ spots are already disposed to their final form when they emerge. The driving force for the shear is spot motion, either flux emergence or the forward motion of p spots in an inverted magnetic configuration. We observe the following phenomena preceding great flares:

1. δ spots, preferentially Types 1 and 2.
2. Umbrae obscured by Hα emission.
3. Bright Hα emission marking flux emergence and reconnection.
4. Greatly sheared magnetic configurations, marked by penumbral and Hα fibrils parallel to the inversion line.

We assert that with adequate spatial resolution one may predict the occurrence of great flares with these indicators.     

Motions and magnetic fields in filaments

Based on the developed method of jointly using data on the magnetic fields and brightness of filaments and coronal holes (CHs) at various heights in the solar atmosphere as well as on the velocities in the photosphere, we have obtained the following results:

The upward motion of matter is typical of filament channels in the form of bright stripes that often surround the filaments when observed in the HeI 1083 nm line.

The filament channels observed simultaneously in Hα and HeI 1083 nm differ in size, emission characteristics, and other parameters. We conclude that by simultaneously investigating the filament channels in two spectral ranges, we can make progress in understanding the physics of their formation and evolution.

Most of the filaments observed in the HeI 1083 nm line consist of dark knots with different velocity distributions in them. A possible interpretation of these knots is offered.

The height of the small-scale magnetic field distribution near the individual dark knots of filaments in the solar atmosphere varies between 3000 and 20000 km.

The zero surface separating the large-scale magnetic field structures in the corona and calculated in the potential approximation changes the inclination to the solar surface with height and is displaced in one or two days.

The observed formation of a filament in a CH was accompanied by a significant magnetic field variation in the CH region at heights from 0 to 30000 km up to the change of the predominant field sign over the entire CH area. We assume that this occurs at the stage of CH disappearance.

     

Spectral analyses of solar photospheric fluctuations

Successful subtraction of instrumental background variations has permitted spectral analyses of two-dimensional measurement arrays of granulation brightness fluctuations at the center of the disk, arrays obtained from Stratoscope I, 1959B-flight, high-resolution frames B1551 and B3241.

1. RMS's, uncorrected for instrumental blurring, are 0.0850 of mean intensity for B1551 and 0.0736 for B3241, somewhat higher than other determinations. These between-frame and between-investigation differences probably result from a combination of calibration errors, frame resolution differences, and, most likely, granulation pattern differences.
2. Significant variations over each array of mean intensities and RMS's, determined for sub-arrays with dimensions in the 2500–10000 km range, indicate spatial brightness and RMS variations larger than the ‘scale’ of the granulation pattern, supporting a turbulent interpretation of photospheric convection.
3. One-dimensional power-spectra shapes provide objective and discriminating criteria for determining granulation pattern differences and, possibly, frame resolution.
4. Two-dimensional power spectra show small, essentially random deviations from axial symmetry which lie almost entirely within the 50% confidence limits.
5. Spectral densities and fluctuation power spectra, computed from the two-dimensional power spectra and corrected for instrumental blurring, noise, and blemishes, have a useable radial wavenumber range nearly double that of earlier Stratoscope I analyses.
6. Corrected RMS's obtained from the corrected fluctuation power spectra, 0.145 ± 0.046 for B1551 and 0.136 ± 0.048 for B3241, depend critically on the accuracy of the correction.
7. The spectra's wavenumber range includes the granulation-fluctuation-producing domain but not the Kolmogoroff domain of turbulence spectra.

     

Energy changes of cosmic rays in the interplanetary region

A clarification and discussion of the energy changes experienced by cosmic rays in the interplanetary region is presented. It is shown that the mean time rate of change of momentum of cosmic rays reckoned for a fixed volume in a reference frame fixed in the solar system is 〈p〉 =p V·G/3 (p=momentum,V is the solar wind velocity andG=cosmic-ray density gradient). This result is obtained in three ways:

1. by a rearrangement and reinterpretation of the cosmic-ray continuity equation;
2. by using a scattering analysis based on that of Gleeson and Axford (1967);
3. by using a special scattering model in which cosmic-rays are trapped in ‘magnetic boxes’ moving with the solar wind.

The third method also gives the rate of change of momentum of particles within a moving ‘magnetic box’ as 〈p〉_ad = ?p ?·V/3, which is the adiabatic deceleration rate of Parker (1965). We conclude that ‘turnaround’ energy change effects previously considered separately are already included in the equation of transport for cosmic rays.     

Flows around sunspots and pores

We report on three sequences of high-resolution white-light and magnetogram observations obtained in the summer of 1989. The duration of sub-arcsecond seeing was three to four hours on each day. Study of the white-light and magnetogram data yields the following results:

1. For all but one of the sunspots we have observed, both dark fibrils and bright grains in the inner part of the penumbra of sunspots move toward the umbra with a speed of about 0.5 km s^-1. In the outer part of the penumbra, movement is away from the umbra. The one exception is a newly formed spot, which has inflow only in its penumbra.
2. Granular flows converge toward almost every pore, even before its formation. Pores are observed to form by the concentration of magnetic flux already existing in the photosphere. The pores (or small sunspots), in turn, then move and concentrate to form bigger sunspot.
3. We followed an emerging flux region (EFR) from 29 to 31 July, 1989 that was composed of a large number of bipoles with magnetic polarities mixed over a large area in the first day of its birth. As time went on, polarities sorted out: the leading polarity elements moved in one direction; the following, the opposite. During the process a large number of cancellations occurred, with some sub-flares and surges observed simultaneously. After about 24 hours, the positive and negative fluxes were essentially separated.
4. We find two kinds of photospheric dark alignments in the region of new flux emergence: (a) alignments connecting two poles of opposite magnetic polarity form the tops of rising flux tubes; (b) alignments corresponding to the magnetic flux of one polarity, which we call elongated pores.

     

Flares and changing magnetic fields

An observational study of maps of the longitudinal component of the photospheric fields in flaring active regions leads to the following conclusions:

1. The broad-wing Hα kernels characteristic of the impulsive phase of flares occur within 10″ of neutral lines encircling features of isolated magnetic polarity (‘satellite sunspots’).
2. Photospheric field changes intimately associated with several importance 1 flares and one importance 2B flare are confined to satellite sunspots, which are small (10″ diam). They often correspond to spot pores in white-light photographs.
3. The field at these features appears to strengthen in the half hour just before the flares. During the flares the growth is reversed, the field drops and then recovers to its previous level.
4. The magnetic flux through flare-associated features changes by about 4 × 10¹⁹ Mx in a day. The features are the same as the ‘Structures Magnétiques Evolutives’ of Martres et al. (1968a).
5. An upper limit of 10²¹ Mx is set for the total flux change through McMath Regions 10381 and 10385 as the result of the 2B flare of 24 October, 1969.
6. Large spots in the regions investigated did not evince flux changes or large proper motions at flare time.
7. The results are taken to imply that the initial instability of a flare occurs at a neutral point, but the magnetic energy lost cannot yet be related to the total energy of the subsequent flare.
8. No unusual velocities are observed in the photosphere at flare time.

     

Orbital eccentricity study for the spherical component of our galaxy

An analysis of the data concerning high-velocity stars from Eggen's catalogue aimed at a determination of the approximate slope of the mass function for the spherical component of our Galaxy, and at estimating the local circular velocity, as well as the local rotation velocity, as by-products, has been performed. Our conclusions are that:

1. A linear dependence of the mass on the radius is very likely;
2. the value of the limiting radius is most likely equal to (40±10) kpc;
3. the two local velocities are approximately equal to each other, being both equal to (230±30) km s^?1;
4. the local escape velocity appears to be most likely equal to (520±30) km s^?1;
5. the total mass of a corona, obtained in this way, is (5±1)×10¹¹ M _⊙.

     

Observations of moving magnetic features near sunspots

The properties of small (< 2″) moving magnetic features near certain sunspots are studied with several time series of longitudinal magnetograms and Hα filtergrams. We find that the moving magnetic features:

1. Are associated only with decaying sunspots surrounded entirely or in part by a zone without a permanent vertical magnetic field.
2. Appear first at or slightly beyond the outer edge of the parent sunspot regardless of the presence or absence of a penumbra.
3. Move approximately radially outward from sunspots at about 1 km s^?1 until they vanish or reach the network.
4. Appear with both magnetic polarities from sunspots of single polarities but appear with a net flux of the same sign as the parent sunspot.
5. Transport net flux away from the parent sunspots at the same rates as the flux decay of the sunspots.
6. Tend to appear in opposite polarity pairs.
7. Appear to carry a total flux away from sunspots several times larger than the total flux of the sunspots.
8. Produce only a very faint emmission in the core of Hα.

A model to help understand the observations is proposed.     

Drift theory of charged particles in electric and magnetic fields

In this paper we review the drift theory of charged particles in electric and magnetic fields. No new physical interpretations are added to this classical topic, but through an alternative, simplified derivation of the guiding centre velocity, several complexities are eliminated and possible misconceptions of the theory are clarified. It is shown that:

1. The curvature/gradient drift velocity in the magnetic field, averaged over a particle distribution function is to lowest order in the direction of?×B/B ², while the average particle velocity is in the direction ofB×? P withP the scalar particle pressure.
2. These drift directions are correct for first-order expansions of the particle distribution function, and only second-order or higher expansions change these directions.
3. The?×B/B ² drift, which is the standard gradient plus curvature drift, and which is usually considered as a ‘single particle’ drift, need not be ‘reconciled’ with theB×? P, or ‘macroscopic, collective’ drift, as is often asserted in the literature. They are in fact related per definition and we show how.
4. When viewed in fixed momentum intervals (p,p+dp), the so-called Compton-Getting factor enters into the electric field (E×B)/B ² drift term.
5. The results are independent of the scale length of variation ofE andB, in contrast to existing drift theory. We discuss the implications of this result for three important cases.

     

The Alfvén-wave theory of solar flares

Evidence is discussed showing that a representative solar flare event comprises three or more separate but related phenomena requiring separate mechanisms. In particular it is possible to separate the most energetic effect (the interplanetary blast) from the thermal flare and from the rapid acceleration of particles to high energies. The phenomena are related through the magnetic structure characteristic of a composite flare event, being a bipolar surface field with most of its field lines ‘closed’. Of primary importance are helical twists on all scales, starting with the ‘flux rope’ of the spot pair which was fully twisted before it emerged. Subsequent untwisting by the upward propagation of an Alfvén twist wave provides the main flare energy.

1. The interplanetary blast model is based on subsurface, helically twisted flux ropes which erupt to form spots and then transfer their twists and energy by Alfvén-twist waves into the atmospheric magnetic fields. The blast is triggered by the prior-commencing flash phase or by a coronal wave.
2. The thermal flare is explained in terms of Alfvén waves travelling up numerous ‘flux strands’ (Figure 3) which have frayed away from the two flux ropes. The waves originate in interaction (collisions, bending, twisting, rubbing) between subsurface flux strands; the sudden flash is caused by a collision. The classical twin-ribbon flare results from the collision of a flux rope with a tight bunch of S-shaped flux strands.
3. The impulsive acceleration of electrons (hard X-ray, EUV, Hα and radio bursts) is tentatively attributed to magnetic reconnection between fields in two parallel, helically twisted flux strands in the low corona.
4. Flare (Moreton) waves in the corona have the same origin as the interplanetary blast. Sympathetic flares represent only the start of enhanced activity in a flare event already in the slow phase. Filament activation also occurs during the slow phase as twist Alfvén waves store their energy in the atmosphere.
5. Flare ejecta are caused by Alfvén waves moving up flux strands. Surges are attributed to packets of twist Alfvén waves released into bundles of flux strands; the waves become non-linear and drive plasma upwards. Spray-type prominences result from accumulations of Alfvén wave energy in dome-shaped fields; excessive energy density eventually explodes the field.

     

Distribution of active regions on the solar surface estimated by statistical methods

Statistical properties of solar active regions (AR) have been studied. In particular, (1) the distribution of ARs by their areas and importances using normal and lognormal distribution laws; (2) it was checked whether the distribution of the ARs' birth sites satisfies the Poisson distribution law (the so-called ‘law of rare events’). Observational data of 1979–1982 have been used and our conclusions are as follows:

1. As regards the areas, the distribution of the ARs that emerged near or on the borders of the large-scale background fields is normal or lognormal.
2. As regards the importances, the distribution of all ARs is lognormal.
3. The distribution of ARs that emerged far from background field borders is not normal.
4. ARs are not casual or rare events on the Sun.

     

Some properties of velocity fields in the solar photosphere

Photoelectric measurements of Doppler shifts of various Fraunhofer lines obtained with the Capri magnetograph were analysed. The height dependence of the supergranular and oscillatory motions, as well as the two dimensional structure of these velocity fields is investigated. The most interesting results are the following:

1. The oscillatory and supergranular motions are still clearly present in very deep photospheric layers as detected e.g. by means of the Ci line at 5380.3 Å.
2. Whereas the vertical motions (both of oscillation and supergranulation) increase with height, the horizontal component of the supergranular flow is found to be decreasing slightly.
3. Aperiodic horizontal motions are observed in the photospheric layers, which are probably connected with the process of excitation of the oscillatory field.
4. There is no simple way of describing the oscillatory field in terms of independently oscillating ‘cells’, since the two-dimensional pattern changes its appearance drastically already in a fraction of one oscillation period.
5. The correlation obtained by previous observers between vertical stationary motions, the chromospheric network and magnetic fields in particular is confirmed.

     

Correlation and spectral analysis of daily solar radio flux

Correlation and spectral analysis of solar radio flux density and sunspot number near the maximum of the sunspot cycle has indicated the existence of

1. long period amplitude modulation of the slowly varying component (SVC) of radio emission
2. coronal storage over a period of the order of three solar rotations
3. fast decay (one solar rotation period or less) of gyromagnetic emissions from radio sources
4. shift in location of chromospheric sources compared to those of either the upper corona or the photosphere.

     

Magnetic knots near a sunspot

Observations of longitudinal and transversal fields and of radial velocities in the magnetic ‘knots’ close to a sunspot were made with the help of Sayan Observatory magnetograph with spatial resolution 1″.2 x 1″.8. The analysis led to following conclusions:

1. The magnetic field in the knots is mainly vertical. The mean inclination of the magnetic-field vector to the vertical direction is equal to 26°.
2. The phenomenon of darkening is connected with essentially vertical fields and brightening in the faculae with the horizontal fields on the sun.
3. An inverse relation between the value of darkening and the inclination of the field vector to the vertical direction and a direct relation on the longitudinal magnetic-field strength exist for the magnetic knots.
4. The magnetic knots in the active region are located in the Hα flocculi near the line where the radial velocity is changing sign in the photosphere.