On the relative and absolute ages of seven lunar front face basins: I. From Viscosity Arguments

The seven basins, Orientale, Imbrium, Crisium, Nectaris, Humorum, and an unnamed basin between Werner and the Altai Ring show rims whose absolute and relative heights are correlated with the sharpness and crispness of the features. On the assumption that the decline in average outer rim height, not scarp height, measures the age of the basin and also that the decline represents a hot creep of rocks of very high viscosity, absolute ages were derived. Basins were found to increase in age in the sequence listed above, with a range from about 3.82 to 4.30 × 10⁹ years. The average or effective viscosity of the surface layers down to whatever level was involved in the creep was calculated as increasing from 9.46 × 10²⁴ poises at about 4.30 × 10⁹ years to 1.86 × 10³⁰ poises at present.It should be clear at the outset what the assumptions and associated observations are and why they are necessary to a solution. They will be listed in this abstract and expanded upon in the text.

• 1.(1)The original rim height of each basin was a function only of basin diameter.
• 2.(2)The original rim height was given by Pike's (1983) relation for fresh craters extrapolated to basin diameters.
• 3.(3)The present rim height is that of the most prominent ring structure.
• 4.(4)The smaller rim height of all seven basins, relative to the height predicted by (2) is due largely to creep in the lunar rocks down to some undetermined level. Other forces may contribute to the sinking of the rims, but these are considered to be of lesser importance and are discussed in the text.
• 5.(5)The relative ages of the seven basins are as given in Table I. This sequence differs slightly from that of Wilhelms (1984), for example, but it is that found in Baldwin, 1974, Baldwin, 1987 and is consistent with the results of this paper.
• 6.(6)The age of formation of the Imbrium basin (3.85 × 10⁹ years) inferred from lunar sample studies (particularly Apollo 15) is correct.
• 7.(7)The age of formation of the Serenitatis basin (3.87 +/− .04) × 10⁹ years, inferred from petrologic and geochemical studies of Apollo 17 boulders is incorrect. This is not an assumption, but is a result of the analysis of this paper.
• 8.(8)The rheology of the Moon may be described, for the purposes of this paper, by an effective viscosity valid throughout the layers involved in the creep.
• 9.(9)This effective viscosity is used as a tool to determine basin ages and is not important in itself. It does appear to vary in the same range as terrestrial rocks, but not the lithosphere.
• 10.(10)Other factors such as isostasy, shaking due to jar from later impacts, modification due to rim relief by ejecta, and erosion from small impacts are all close to exponential in nature, declining toward the present, and hence may be included in the determination of the effective viscosity.
• 11.(11)The rim height of the Imbrium basin subsided by 25 m in the last 2.5 × 10⁹ years. This value was chosen arbitrarily. It could have been 250–300 m and the basin ages would not have been affected except for Orientale and that only minutely.
• 12.(12)The effective viscosity of the Moon was observed to change continuously and monotonically with time.
• 13.(13)Judging by Table III, the probable error of an absolute age is in the range of 10 to 50 × 10⁶ years. It is difficult to determine exactly what this means. It will be constrained by points (14), (15), and (16).
• 14.(14)An error in the age determination should not be large enough to alter the relative ages of the basins, judging by crater counts (Baldwin, 1974, Baldwin, 1987).
• 15.(15)If the viscosity of the Moon declined in the post-Imbrium period of mare formation the only basin to be affected would be Orientale and this by no great amount inasmuch as the basin is older than nearby maria.
• 16.(16)If the effective viscosity were less than about 10²² poises at the time of the oldest basin then presaturation surfaces would not show the numerous craters and portions of craters that are obvious in this time span.
• 17.(17)Considerably prior to the time of saturation the outer layers of the Moon had a low enough viscosity so that they could not retain the record of the then occuring cratering.
• 18.(18)The approximations of this paper were adopted because it does not appear possible to make an unambiguous selection from the more elegant mathematical treatments of creep and isostasy that would lead toward reasonable ages for the giant basins.

     

Long-term prediction of solar activity

After reviewing the various methods used so far in the long-term prediction of solar activity, we have opted to use methods based on the 80-year period and probability considerations. Our predictions for Cycle No. 21 are:

• 1.Maximum of the annual average Wolf number: 54 – 94
• 2.Minimum of the annual average Wolf number: 0 – 4
• 3.Epoch of maximum: first half of 1981 - first half of 1983
• 4.Epoch of minimum: first half of 1976 - first half of 1977

A comparison of our predictions with others is made.     

The role of spiral sunspots in the solar activity of October 1972

From a collection of photographs of 192 knots relating to 68 flares of the activity region McMath 12094 of October 1972, and of spots and prominences, a correlation analysis was made which led to the following conclusions:

• 1.1. Flare activity began in the region above the leader spot A(N) which had the highest twisted magnetic energy of all the spots in the group. It rapidly developed to the region above the follower S-spots, and then to regions outside the spot area. This evolutionary process was in step with the morphological changes and the rotational movement of the spots in the group.
• 2.2. The flare activity around Spot A began in the interior of the twisted magnetic tube and was in step with the untwisting of magnetic features, and the flare knots spread from the vicinity of the umbra to the outside of the penumbra. When the knotted structure completely untwisted itself, the overall flare activity decreased in the region.
• 3.3. A spiral spot was found to be a strong centre of attraction below a giant coronal prominence. This shows that coronal prominences are formed through the contraction of strong electrical currents in the corona.

We also discussed the question whether the twisted magnetic structure of spot A could have provided most of the energy in the flare activity and lead to an instability.     

Stochastic theory of turbulent convection in pulsating variables

The dynamic behaviour of turbulent convection is studied. We have solved the dynamic equations of the correlation functions and obtained expressions for the steady and pulsating components. The main results are:

• 1.(1) Vitense's and Öpik's phenomenological theories can be recovered as special cases of the steady component of the present statistical theory.
• 2.(2) The dynamic process of turbulent convection decreases the viscous forces and the pulsating amplitudes of heat convection and turbulent pressure, and gives rise to relative lags in phase.
• 3.(3) A stability analysis indicates the following: a) The direct effect of the energy convection (thermal and mechanical) is to promote pulsational instability below the ionising region of the abundant elements, while in the region itself, the energy convection may produce a damping effect. b) Turbulent viscosity always produces a damping effect. In the quasi-adiabatic region, turbulent pressure produces a destabilising effect. c) The energy transfer by molecular stresses (gas pressure and molecular viscosity) acts in the opposite sense to the turbulent Reynolds stress. These effects could account for the location of the red edge of the pulsational instability strip on the HR diagram and for the light variability of the red giants and supergiants.

     

Outstanding aeronomy problems at Venus

Of all the non-terrestrial ionospheres and thermospheres in our solar system those of Venus have been explored and studied the most. This is mainly because of the 14 year exploration of the well-instrumented Pioneer Venus spacecraft and the theoretical studies prompted by the resulting observational information. However, there are still a number of areas where there are important scientific questions that remain unanswered. These areas include:

• (i)dynamics of the thermosphere,
• (ii)the energy mechanisms/sources responsible for maintaining the elevated plasma temperatures in the ionosphere,
• (iii)airglow/aurora intensities and their sources, and
• (iv)hot atom populations.

Venus Express is likely to help address some of the questions related to the areas listed under (i), (iii) and (iv) above.     

The chemical evolution of the solar neighbourhood: the effect of binaries

In this paper we compute the time evolution of the elements (⁴He, ¹²C, ¹⁴N, ¹⁶O, ²⁰Ne, ²⁴Mg, ²⁸Si, ³²S, ⁴⁰Ca and ⁵⁶Fe) and of the supernova rates in the solar neighbourhood by means of a galactic chemical evolutionary code that includes in detail the evolution of both single and binary stars. Special attention is payed to the formation of black holes.Our main conclusions:

• •in order to predict the galactic time evolution of the different types of supernovae, it is essential to compute in detail the evolution of the binary population,
• •the observed time evolution of carbon is better reproduced by a galactic model where the effect is included of a significant fraction of intermediate mass binaries,
• •massive binary mass exchange provides a possible solution for the production of primary nitrogen during the very early phases of galactic evolution,
• •chemical evolutionary models with binaries or without binaries but with a detailed treatment of the SN Ia progenitors predict very similar age–metallicity relations and very similar G-dwarf distributions whereas the evolution of the yields as function of time of the elements ⁴He, ¹⁶O, ²⁰Ne, ²⁴Mg, ²⁸Si, ³²S and ⁴⁰Ca differ by no more than a factor of two or three,
• •the observed time evolution of oxygen is best reproduced when most of the oxygen produced during core helium burning in ALL massive stars serves to enrich the interstellar medium. This can be used as indirect evidence that (massive) black hole formation in single stars and binary components is always preceded by a supernova explosion.

     

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.     

Problems with non-thermal models for the narrow line gamma rays reported from SS 433

The jet/grain model proposed by Ramatyet al. (1984, hereafter abbreviated as RKL) for production of the narrow gamma-ray lines reported from SS433 is examined and shown to be untenable on numerous grounds. Most importantly:

1. The huge Coulomb collisional losses (W _c?2×10⁴¹ erg s^?1) from the jet, which would necessarily accompany non-thermal production of the gamma rays, demands a jet acceleration/collimation process acting over a very long range and with a power at least 10² times the Eddington limit for any stellar object.
2. There is a collisional thick target limit (irrespective of jet mass) to the gamma ray yield per interstellar proton. Consequently, the gamma-ray data demand an improbably high interstellar density (?10⁹ cm^?3).
3. For the grains to be kept cool enough (?3000 K) to survive the heating rateW _c either by radiation or jet expansion would demand a ‘jet’ wider than its length and so inconsistent with narrow lines. In the case of radiative cooling, the resultant IR flux would exceed the observed values by a factor ?10⁴.
4. Light scattered on the jet grain mass required would be highly polarized, contrary to observations, unless the jet was optically thick to grains, again precluding their radiative cooling.
5. To avoid unacceptable precessional broadening of the gamma-ray lines demands an emitting jet length ?0.5 days atv=0.26c. This increases the necessary mass loss rate by a factor ?10 over the values obtained by RKL who assumed a 4-day ‘flare’.
6. The model also predicts rest energy gamma-ray lines which are not observed.

     

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.

     

Radio science investigations by VeRa onboard the Venus Express spacecraft

The Venus Express Radio Science Experiment (VeRa) uses radio signals at wavelengths of 3.6 and 13 cm (“X”- and “S”-band, respectively) to investigate the Venus surface, neutral atmosphere, ionosphere, and gravity field, as well as the interplanetary medium. An ultrastable oscillator (USO) provides a high quality onboard reference frequency source; instrumentation on Earth is used to record amplitude, phase, propagation time, and polarization of the received signals. Simultaneous, coherent measurements at the two wavelengths allow separation of dispersive media effects from classical Doppler shift.VeRa science objectives include the following:

• (1)Determination of neutral atmospheric structure from the cloud deck (approximately 40 km altitude) to 100 km altitude from vertical profiles of neutral mass density, temperature, and pressure as a function of local time and season. Within the atmospheric structure, search for, and if detected, study of the vertical structure of localized buoyancy waves, and the presence and properties of planetary waves.
• (2)Study of the H₂SO₄ vapor absorbing layer in the atmosphere by variations in signal intensity and application of this information to tracing atmospheric motions. Scintillation effects caused by radio wave diffraction within the atmosphere can also provide information on small-scale atmospheric turbulence.
• (3)Investigation of ionospheric structure from approximately 80 km to the ionopause (<600 km), allowing study of the interaction between solar wind plasma and the Venus atmosphere.
• (4)Observation of forward-scattered surface echoes obliquely reflected from selected high-elevation targets with anomalous radar properties (such as Maxwell Montes). More generally, such bistatic radar measurements provide information on the roughness and density of the surface material on scales of centimeters to meters.
• (5)Detection of gravity anomalies, thereby providing insight into the properties of the Venus crust and lithosphere.
• (6)Measurement of the Doppler shift, propagation time, and frequency fluctuations along the interplanetary ray path, especially during periods of superior conjunction, thus enabling investigation of dynamical processes in the solar corona.

     

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.     

Exotic ion H 3 ++ in strong magnetic fields

1. The exotic system H ₃ ⁺⁺ (which does not exist without magnetic field) exists in strong magnetic fields:
   1. In triangular configuration for B≈10⁸–10¹¹?G (under specific external conditions)
   2. In linear configuration for B>10¹⁰?G
2. In the linear configuration the positive z-parity states 1σ _g , 1π _u , 1δ _g are bound states
3. In the linear configuration the negative z-parity states 1σ _u , 1π _g , 1δ _u are repulsive states
4. The H ₃ ⁺⁺ molecular ion is the most bound one-electron system made from protons at B>3×10¹³?G

Possible application: The H ₃ ⁺⁺ molecular ion may appear as a component of a neutron star atmosphere under a strong surface magnetic field B=10¹²–10¹³?G.     

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.

     

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.

     

A morphological study of a solar active region in August 1972: II. A correlation analysis of morphologies of flares and spots

A correlation analysis of the morphology of nine solar flares of a large solar active region in August 1972 and the morphology of the fine structure of sunspot group in this region has led to the following conclusions:

• 1.1. There is a certain correlation between the outburst of the first large flare, occurring at 0355 UT. of August 2nd and the morphological variations of photospheric spots both in time and in spatial positions.
• 2.2. All the preliminary bright points of the nine flares on the both sides of the filament and their main morphological development are also closely correlated with the spiral structure of the spots “O” and “B”.
• 3.3. All the directions of the bright ribbons of flares and filaments (consisting of paralled fibrils) in the chromosphere, the serpent-like long fibres of penumbra on the east of the spot “O” and the line H_∥ = 0 in the photosphere are consistent with one another. This can be regarded as a morphological evidence in favour the opinion that the outburst of flares is propagated along the horizontal magnetic field on the surface of the sun.

     

A varying cosmological term as a link between cosmology and microphysics

The Weinberg relation (which connects the Hubble constantH to the mass of a typical elementary particle) is an empirical relation hitherto unexplained. I suggest an explanation based on the Zel'dovich energy tensor of vacuum in a Robertson-Walker universe with constant deceleration parameter,q = const. This model leads to

1. the Weinberg relation,
2. a varying cosmological term Λ scaling asH ²,
3. a varying gravitational constantG scaling asH,
4. a matter creation process throughout the universe at the rate 10^?47 g s^?1 cm³,
5. a deceleration parameter in the range -1 to 1/2, which allows a horizon-free universe and makes the lawG/H = constant, consistent with the Viking lander data on the orbit of planet Mars.

     

Observation of the impulsive phase of a simple flare

We present a broad range of complementary observations of the onset and impulsive phase of a fairly large (1B, M1.2) but simple two-ribbon flare. The observations consist of hard X-ray flux measured by the SMM HXRBS, high-sensitivity measurements of microwave flux at 22 GHz from Itapetinga Radio Observatory, sequences of spectroheliograms in UV emission lines from Ov (T ≈ 2 × 10⁵ K) and Fexxi (T ≈ 1 × 10⁷ K) from the SMM UVSP, Hα and Hei D₃ cine-filtergrams from Big Bear Solar Observatory, and a magnetogram of the flare region from the MSFC Solar Observatory. From these data we conclude:

1. The overall magnetic field configuration in which the flare occurred was a fairly simple, closed arch containing nonpotential substructure.
2. The flare occurred spontaneously within the arch; it was not triggered by emerging magnetic flux.
3. The impulsive energy release occurred in two major spikes. The second spike took place within the flare arch heated in the first spike, but was concentrated on a different subset of field lines. The ratio of Ov emission to hard X-ray emission decreased by at least a factor of 2 from the first spike to the second, probably because the plasma density in the flare arch had increased by chromospheric evaporation.
4. The impulsive energy release most likely occurred in the upper part of the arch; it had three immediate products:

1. An increase in the plasma pressure throughout the flare arch of at least a factor of 10. This is required because the Fexxi emission was confined to the feet of the flare arch for at least the first minute of the impulsive phase.
2. Nonthermal energetic (～ 25 keV) electrons which impacted the feet of the arch to produce the hard X-ray burst and impulsive brightening in Ov and D₃. The evidence for this is the simultaneity, within ± 2 s, of the peak Ov and hard X-ray emissions.
3. Another population of high-energy (～100keV) electrons (decoupled from the population that produced the hard X-rays) that produced the impulsive microwave emission at 22 GHz. This conclusion is drawn because the microwave peak was 6 ± 3 s later than the hard X-ray peak.

     

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.

     

Microwave emission above steady and moving sunspots

Two-dimensional maps of radio brightness temperature and polarization, computed assuming thermal emission with free-free and gyroresonance absorption, are compared with observations of active region 2502, performed at Westerbork at λ = 6.16 cm during a period of 3 days in June 1980. The computation is done assuming a homogeneous model in the whole field of view (5′ × 5′) and a force-free extrapolation of the photospheric magnetic field observed at MSFC with a resolution of 2″.34. The mean results are the following:

1. A very good agreement is found above the large leading sunspot of the group, assuming a potential extrapolation of the magnetic field and a constant conductive flux in the transition region ranging from 2 × 10⁶ to 10⁷ erg cm^?2s^?1.
2. A strong radio source, associated with a new-born moving sunspot, cannot be ascribed to thermal emission. It is suggested that this source may be due to synchrotron radiation by mildly relativistic electrons accelerated by resistive instabilities occurring in the evolving magnetic configuration. An order-of-magnitude computation of the expected number of accelerated particles seems to confirm this hypothesis.

     

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.