Global solar activity on long time scales

We report the results of the application of our approach to study the behavior of solar activity in the past, where:

When reconstructing the variations of solar activity, geomagnetic parameters, and the interplanetary magnetic field in the past we select a sequence of increasing time scales, which can be naturally represented by the potentials of available observational data. We select a total of four time scales: 150–200 years, 400 years, 1000 years, and 10000 years.

When constructing the series of each successive (in terms of length) time scale we use the data of the previous time scale as reference data.

We abandon, where possible, the series of traditional statistical parameters in favor of the series of physical parameters.

When deriving the relations between any parameters of solar activity, geomagnetic disturbance, and the interplanetary magnetic field, we take into account the differential nature of relations on different time scales. To this end, we use the earlier proposed MSR and DPS methods.

To verify the resulting reconstructions, we use the “principle of witnesses”, which uses independent (in some cases, indirect) information as initial data.

     

Keynote address: Outstanding problems in solar physics

Celebrating the diamond jubilee of the Physics Research Laboratory (PRL) in Ahmedabad, India, we look back over the last six decades in solar physics and contemplate on the ten outstanding problems (or research foci) in solar physics:

1. The solar neutrino problem
2. Structure of the solar interior (helioseismology)
3. The solar magnetic field (dynamo, solar cycle, corona)
4. Hydrodynamics of coronal loops
5. MHD oscillations and waves (coronal seismology)
6. The coronal heating problem
7. Self-organized criticality (from nanoflares to giant flares)
8. Magnetic reconnection processes
9. Particle acceleration processes
10. Coronal mass ejections and coronal dimming

The first two problems have been largely solved recently, while the other eight selected problems are still pending a final solution, and thus remain persistent Challenges for Solar Cycle 24, the theme of this jubilee conference.     

The results of statistical analysis of the coronal profiles above the solar active regions

In order to establish some regularities or variations in the distribution of widths and intensities of the coronal line profiles λ 5303 and λ 6374 depending upon the solar activity, a statistical analysis was made for more than 3000 profiles (the data covering the period 1966–1972). The following results obtained:

1. The distribution of coronal line profile widths changes depending upon the solar activity phase.
2. The character of the relation between the intensities and widths varies with variation of the solar activity phase.

     

Solar flare physics enlivened by TRACE and RHESSI

The Transition Region and Coronal Explorer (TRACE) gave us the highest EUV spatial resolution and the Ramaty High Energy Solar Spectrometric Imager (RHESSI) gave us the highest hard X-ray and gammaray spectral resolution to study solar flares. We review a number of recent highlights obtained from both missions that either enhance or challenge our physical understanding of solar flares, such as:

1. Multi-thermal Diagnostic of 6.7 and 8.0 keV Fe and Ni lines
2. Multi-thermal Conduction Cooling Delays
3. Chromospheric Altitude of Hard X-Ray Emission
4. Evidence for Dipolar Reconnection Current Sheets
5. Footpoint Motion and Reconnection Rate
6. Evidence for Tripolar Magnetic Reconnection
7. Displaced Electron and Ion Acceleration Sources.

     

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.

     

Ephemeral active regions in 1970 and 1973

A study of ephemeral active regions (ER) identified on good quality full-disk magnetograms reveals:

1. On the average 373 and 179 ER were present on the Sun in 1970 and 1973 respectively. The number varies with the solar cycle.
2. The median lifetime of ER depends on observation quality and selection rules but is estimated as about 12 hr for our data.
3. The latitude distribution is very broad but not uniform. The distribution peaks near the equator and shows variations similar to distributions of large active regions.
4. The longitude distribution is essentially homogeneous.
5. The spatial orientation of ER is almost random. In 1973 there is a hint of an excess of new cycle orientations at high latitudes.

A comparison of parameters of ER and regular active regions suggests that ER are the small-scale end of a broad spectrum of active regions. The role of ER in the light of present theories of solar activity is investigated but is not yet clear. Heating of the chromosphere and corona may be significantly affected by ER.     

On the latitude migration of polar faculae in the solar activity cycle period: 1970–1978

Coordinates of polar faculae have been measured and processed using daily photoheliograms of the Kislovodsk Station of the Pulkovo observatory with the final goal of studying their latitude distribution during the solar cycles 20–21. The results obtained are as follows:

1. The first polar faculae emerge immediately after the polarity inversion of the solar magnetic field at the latitudes from 40° to 70° with the average ?-55°.
2. The zone of the emergence of polar faculae migrates poleward during the period between the neighbouring polarity inversions of the solar magnetic field. This migration is about 20° for 8 years, which corresponds to a velocity of 0.5 m s^-1.
3. The maximum number of polar faculae was reached at the activity minimum (1975–1976).
4. The last polar faculae were observed in the second half of 1978 at the latitudes from 70° to 80°.

     

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.     

An evaluation of the possibility of studying flare plasma turbulence using the satellites of Hei line forbidden components

Using the Baranger-Mozer method, we explore the possibility of diagnosing the flare plasma of forbidden Hei lines, that permits the determination of the plasma oscillation frequency and noise level. Examination of the Hei lines observed in solar flare has led us to conclude that:

1. the appearance of satellites of forbidden components in the flares spectrum, due to turbulent electric fields, is the most probable for Hei 3819.606 Å lines;
2. the Baranger-Mozer method is more sensitive to the high-frequency component of turbulent fields than to the low-frequency ones;
3. the upper limit of the turbulent oscillation level in flares is evaluated.

In the spectrum of the solar flare of 26 September, 1963 we detected satellites of the forbidden component of the 3820 Å line and used its relative intensity to derive the level of low-frequency oscillations (～1.5 kVcm^-1).     

Evolution of the solar Lyman alpha flux during four consecutive years

Four consecutive years of a quasi-continuous survey of the solar Lα line are presented. Absolute calibration and aging correction are evaluated producing higher quality measurements which are:

-the total Lα flux,

-the central Lα flux,

-the blue wing flux at 0.33 Å from the center,

-the slope of the blue wing at the same location.

Empirical laws are deduced from this large amount of data giving a relation between these different parameters and the flux integrated over the whole line. Furthermore, other empirical laws are obtained between the total Lα flux and two solar activity indices. These relations give a possibility of evaluating the solar Lα flux even when no observation is available and, as previously found by Prinz (1974), show that to a first approximation the solar Lα flux is composed of a quiet and of an active component. The active component changes with the 27 days period; the quiet one with the 11 yr solar cycle.     

On some aspects of the albedo effect application to the ERS-1 satellite

Some aspects of the perturbative influence of radiation reflected by the Earth's surface on the motion of an artificial satellite are discussed. We concentrate on consequences of the extreme models with anisotropic reflection on the Earth's surface (specular reflection, clouds with anisotropic phase function). The possible effects of Lála's modification of the Earth's albedo nominal value are investigated. The role of the satellite surface optical properties is pointed out in the context of the albedo effect. All mentioned models are purely numerical. The whole message of the paper can be summarized in the following items

-It is very doubtful that the 10^?8 ÷ 10^?9 m s^?2 level is reached when determining the perturbing accelerations caused by the albedo effect in the case of the ERS-1 satellite due to poorly defined optical characteristics of the Earth's atmosphere, the Earth and the satellite's surface.

-In the general case this albedo effect uncertainty level is about 50% with respect to the averaged values, and probably as high as 100% with respect to the instantaneous values of the perturbing accelerations.

     

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.

     

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.

     

Structure and evolutionary history of the solar system,I

1. Introduction and Survey. The method for studying the structure and evolution of the solar system is discussed. It is pointed out that theories that account for the origin of planets alone are basically insufficient. Instead one ought to aim for a general theory for the formation of secondary bodies around a central body, applicable both to planet and satellite formation. A satisfactory theory should not start from assumed properties of the primitive Sun, which is a very speculative subject, but should be based on an analysis of present conditions and a successive reconstruction of the past states.
2. Orbits of Planets and Satellites. As a foundation for the subsequent analysis, the relevant properties of planets and satellites are presented.
3. The Small Bodies. The motion of small bodies is influenced by non-gravitational forces. Collisions (viscosity) are of special importance for the evolution of the orbits. It is pointed out that the focusing property of a gravitational field (which has usually been neglected) leads to the formation of jet streams. The importance of this concept for the understanding of the comet-meteoroid relations and the structure of the asteroidal belt is shown.
4. Resonance Structure. A survey is given of the resonances in the solar system and their possible explanation. It is concluded that in many cases the resonances must already be produced at the times when the bodies formed. It is shown that resonance effects put narrow limits on the post-accretional changes of orbits.
5. Spin and Tides. Tidal effects on planetary spins and satellite orbits are discussed. It is very doubtful if any satellite except the Moon and possibly Triton has had its orbit changed appreciably by tidal effects. The isochronism of planetary and asteroidal spins is discussed, as well as its bearing on the accretional process.
6. Post-accretional Changes in the Solar System. The stability of the solar system and upper limits for changes in orbital and spin data are examined. It is concluded that much of the present dynamic structure has direct relevance to the primordial processes.

     

Rotational Tomography For 3d Reconstruction Of The White-Light And Euv Corona In The Post-Soho Era

Improving our understanding of the mechanisms that energize the solar wind and heat structures in the solar corona requires the development of empirical methods that can determine the three-dimensional (3D) temperature and density distributions with as much spatial and temporal resolution as possible. This paper reviews the solar rotational tomography (SRT) methods that will be used for 3D reconstruction of the solar corona from data obtained by the next generation of space-based missions such as the Solar and Terrestrial Relations Observatory (STEREO), Solar-B and the Solar Dynamics Observatory (SDO). In the next decade, SRT will undergo rapid advancement on several frontiers of 3D image reconstruction:

1. Electron density reconstruction from white-light coronagraph images.
2. Differential emission measure (DEM) reconstruction from EUV images.
3. Dual-spacecraft (STEREO) observing geometry.
4. Fusion of data from multiple spacecraft with differing instrumentation.
5. Time-dependent estimation methods.

Although the principles described apply to many different wavelength regimes, this paper concentrates on white-light and EUV data. Previous work on all of these subjects is reviewed, and major technical issues and future directions are discussed.     

Alfvén waves in the solar atmosphere

We examine the propagation of Alfvén waves in the solar atmosphere. The principal theoretical virtues of this work are: (i) The full wave equation is solved without recourse to the small-wavelength eikonal approximation (ii) The background solar atmosphere is realistic, consisting of an HSRA/VAL representation of the photosphere and chromosphere, a 200 km thick transition region, a model for the upper transition region below a coronal hole (provided by R. Munro), and the Munro-Jackson model of a polar coronal hole. The principal results are:

1. If the wave source is taken to be near the top of the convection zone, where n _H = 5.2 × 10¹⁶ cm^?3, and if B _⊙ = 10.5 G, then the wave Poynting flux exhibits a series of strong resonant peaks at periods downwards from 1.6 hr. The resonant frequencies are in the ratios of the zeroes of J ₀, but depend on B _⊙, and on the density and scale height at the wave source. The longest period peaks may be the most important, because they are nearest to the supergranular periods and to the observed periods near 1 AU, and because they are the broadest in frequency.
2. The Poynting flux in the resonant peaks can be large enough, i.e. P _⊙ ≈ 10⁴–10⁵ erg cm^?2s^?1, to strongly affect the solar wind.
3. ¦δv¦ and ¦δB¦ also display resonant peaks.
4. In the chromosphere and low corona, ¦δv ≈ 7–25 kms^?1 and ¦δB¦ ≈0.3–1.0 G if P _⊙≈10⁴-10⁵ erg cm^?2s^?1.
5. The dependences of ¦δv¦ and ¦δB¦ on height are reduced by finite wavelength effects, except near the wave source where they are enhanced.
6. Near the base, ¦δB¦ ≈ 350–1200 G if P _⊙ ～- 10⁴–10⁵. This means that nonlinear effects may be important, and that some density and vertical velocity fluctuations may be associated with the Alfvén waves.
7. Below the low corona most wave energy is kinetic, except near the base where it becomes mostly magnetic at the resonances.
8. ?₀ < δv ² > v _A or < δB ² > v _A/4π are not good estimators of the energy flux.
9. The Alfvén wave pressure tensor will be important in the transition region only if the magnetic field diverges rapidly. But the Alfvén wave pressure can be important in the coronal hole.

     

Effects of temperature and velocity gradients on doppler widths

We have investigated how the gradients of temperature and expansion velocities will change the emergent profiles from an extended medium in spherical symmetry. Variation of the source function and expansion velocities are assumed. The following variations of temperature are employed:

1. T(r) ; T₀ (isothermal case)
2. T(r) ; T₀(r/r₀)^1/2
3. T(r) ; T₀(r/r₀)^-1
4. T(r) ; T₀(r/r₀)^-2
5. T(r) ; T₀(r/r₀)^-3

The profiles calculated present an interesting feature of broadening.     

The effect of solar parameters on seasonal variation of 5893 Å line intensity at Calcutta

The purpose of this paper is to present the correlation of seasonal variation of 5893 Å line intensity with relative sunspot numbers, solar flare numbers and the variable component of 10.7 cm solar flux. A study has been made and the following important results have been obtained.

1. The intensity of 5893 Å line at Calcutta shows periodic variation with different solar parameters during descending part of secondary peak of 21st solar cycle (1984–1985).
2. 5893 Å line intensity of Mt. Abu also shows periodic variation with solar parameters during the period 1965–1968 when there was a peak phase of 20th solar cycle.
3. A possible explanation for such type of variation is also presented.

     

Nature of variation of Antarctic O3 depletion and its correlation with solar UV radiation

The purpose of this paper is to study the nature of variation of O₃ concentration of Antarctic Survey Stations and its correlation with solar ultraviolet radiation. Solar UV data for the period November 1978 to October 1984 are taken from Solar Geophysical Data Book. In absence of solar UV data for long period, a calibration curve between solar UV radiation and solar flare number (S.F.NO.) is drawn. (A straight line is obtained and correlation coefficient between two variables is 80%). The equation of straight line from least square principle becomes, UV flux = 0.2672 + 2.7578 × 10^?5 × S.F.NO. From this equation UV flux values for long period are calculated from known values of solar flare numbers. O₃ concentration of two Antarctic Survey Stations, Halley Bay (76? S, 27? W) and McMurdo (78? S, 166? E) are considered for analysis and following important results are obtained:

1. Yearly variations of O₃ concentrations and UV radiations are mainly controlled by their October concentrations.
2. Correlation coefficient between O₃ concentration and UV radiation is 62% for the month of October. For the other months it is poor.
3. It is concluded that dramatic decrease of O₃ concentration at Antarctica is independent of solar UV radiation and chemical processes are responsible for special depletion of O₃.

     

Non-thermal processes in large solar flares

We analyze particle acceleration processes in large solar flares, using observations of the August, 1972, series of large events. The energetic particle populations are estimated from the hard X-ray and γ-ray emission, and from direct interplanetary particle observations. The collisional energy losses of these particles are computed as a function of height, assuming that the particles are accelerated high in the solar atmosphere and then precipitate down into denser layers. We compare the computed energy input with the flare energy output in radiation, heating, and mass ejection, and find for large proton event flares that:

1. The ～10–10² keV electrons accelerated during the flash phase constitute the bulk of the total flare energy.
2. The flare can be divided into two regions depending on whether the electron energy input goes into radiation or explosive heating. The computed energy input to the radiative quasi-equilibrium region agrees with the observed flare energy output in optical, UV, and EUV radiation.
3. The electron energy input to the explosive heating region can produce evaporation of the upper chromosphere needed to form the soft X-ray flare plasma.
4. Very intense energetic electron fluxes can provide the energy and mass for interplanetary shock wave by heating the atmospheric gas to energies sufficient to escape the solar gravitational and magnetic fields. The threshold for shock formation appears to be ～10³¹ ergs total energy in >20 keV electrons, and all of the shock energy can be supplied by electrons if their spectrum extends down to 5–10 keV.
5. High energy protons are accelerated later than the 10–10² keV electrons and most of them escape to the interplanetary medium. The energetic protons are not a significant contributor to the energization of flare phenomena. The observations are consistent with shock-wave acceleration of the protons and other nuclei, and also of electrons to relativistic energies.
6. The flare white-light continuum emission is consistent with a model of free-bound transitions in a plasma with strong non-thermal ionization produced in the lower solar chromosphere by energetic electrons. The white-light continuum is inconsistent with models of photospheric heating by the energetic particles. A threshold energy of ～5×10³⁰ ergs in >20 keV electrons is required for detectable white-light emission.

The highly efficient electron energization required in these flares suggests that the flare mechanism consists of rapid dissipation of chromospheric and coronal field-aligned or sheet currents, due to the onset of current-driven Buneman anomalous resistivity. Large proton flares then result when the energy input from accelerated electrons is sufficient to form a shock wave.