共查询到20条相似文献，搜索用时 390 毫秒
1.
David M. Rust 《Solar physics》1972,25(1):141-157
An observational study of maps of the longitudinal component of the photospheric fields in flaring active regions leads to the following conclusions:
- 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’).
- 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.
- 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.
- 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).
- 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.
- Large spots in the regions investigated did not evince flux changes or large proper motions at flare time.
- 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.
- No unusual velocities are observed in the photosphere at flare time.
2.
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:
- 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.
- These drift directions are correct for first-order expansions of the particle distribution function, and only second-order or higher expansions change these directions.
- 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.
- 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.
- 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.
3.
Juan C. López Vieyra Alexander V. Turbiner Nicolais L. Guevara 《Astrophysics and Space Science》2007,308(1-4):493-497
- The exotic system H 3 ++ (which does not exist without magnetic field) exists in strong magnetic fields:
- In triangular configuration for B≈108–1011?G (under specific external conditions)
- In linear configuration for B>1010?G
- In the linear configuration the positive z-parity states 1σ g , 1π u , 1δ g are bound states
- In the linear configuration the negative z-parity states 1σ u , 1π g , 1δ u are repulsive states
- The H 3 ++ molecular ion is the most bound one-electron system made from protons at B>3×1013?G
4.
B. Lehnert 《Astrophysics and Space Science》1977,46(1):61-71
This paper outlines the problems of the quasi-steady matter-antimatter boundary layers discussed in Klein-Alfvén's cosmological theory, and a crude model of the corresponding ambiplasma balance is presented:
- At interstellar particle densities, no well-defined boundary layer can exist in presence of neutral gas, nor can such a layer be sustained in an unmagnetized fully ionized ambiplasma.
- Within the limits of applicability of the present model, sharply defined boundary layers are under certain conditions found to exist in a magnetized ambiplasma. Thus, at beta values less than unity, a steep pressure drop of the low-energy components of matter and antimatter can be balanced by a magnetic field and the electric currents in the ambiplasma.
- The boundary layer thickness is of the order of 2x 0?10/BT 0 1/4 metres, whereB is the magnetic field strength in MKS units andT 0 the characteristic temperature of the low-energy components in the layer.
5.
We define for observational study two subsets of all polar zone filaments, which we call polemost filaments and polar filament bands. The behavior of the mean latitude of both the polemost filaments and the polar filament bands is examined and compared with the evolution of the polar magnetic field over an activity cycle as recently distilled by Howard and LaBonte (1981) from the past 13 years of Mt. Wilson full-disk magnetograms. The magnetic data reveal that the polar magnetic fields are built up and maintained by the episodic arrival of discrete f-polarity regions that originate in active region latitudes and subsequently drift to the poles. After leaving the active-region latitudes, these unipolar f-polarity regions do not spread equatorward even though there is less net flux equatorward; this indicates that the f-polarity regions are carried poleward by a meridional flow, rather than by diffusion. The polar zone filaments are an independent tracer which confirms both the episodic polar field formation and the meridional flow. We find:
- The mean latitude of the polemost filaments tracks the boundary of the polar field cap and undergoes an equatorward dip during each arrival of additional polar field.
- Polar filament bands track the boundary latitudes of the unipolar regions, drifting poleward with the regions at about 10 m s-1.
- The Mt. Wilson magnetic data, combined with a simple model calculation, show that the filament drift expected from diffusion alone would be slower than observed, and in some cases would be equatorward rather than poleward.
- The observation that filaments drift poleward along with the magnetic regions shows that fields of both polarities are carried by the meridional flow, as would be expected, rather than only the f-polarity flux which dominates the strength. This leads to the prediction that in the mid-latitudes during intervals between the passage of f-polarity regions, both polarities are present in nearly equal amounts. This prediction is confirmed by the magnetic data.
6.
Hot spots similar to those in the radio galaxy Cygnus A can be explained by the strong shock produced by a supersonic but classical jet \(\left( {u_{jet}< c/\sqrt 3 } \right)\) . The high integrated radio luminosity (L?2×1044 erg s?1) and the strength of mean magnetic field (B?2×10?4 G) suggest the hot spots are the downstream flow of a very strong shock which generates the ultrarelativistic electrons of energy ?≥20 MeV. The fully-developed subsonic turbulence amplifies the magnetic field of the jet up to 1.6×10?4 G by the dynamo effect. If we assume that the post-shock pressure is dominated by relativistic particles, the ratio between the magnetic energy density to the energy density in relativistic particles is found to be ?2×10?2, showing that the generally accepted hypothesis of equipartition is not valid for hot spots. The current analysis allows the determination of physical parameters inside hot spots. It is found that:
- The velocity of the upstream flow in the frame of reference of the shock isu 1?0.2c. Radio observations indicate that the velocity of separation of hot spots isu sep?0.05c, so that the velocity of the jet isu jet=u 1+u sep?0.25c.
- The density of the thermal electrons inside the hot spot isn 2?5×10?3 e ? cm?3 and the mass ejected per year to power the hot spot is ?4M 0yr?1.
- The relativistic electron density is less than 20% of the thermal electron density inside the hot spot and the spectrum is a power law which continues to energies as low as 30 MeV.
- The energy density of relativistic protons is lower than the energy density of relativistic electrons unlike the situation for cosmic rays in the Galaxy.
7.
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:
- Are associated only with decaying sunspots surrounded entirely or in part by a zone without a permanent vertical magnetic field.
- Appear first at or slightly beyond the outer edge of the parent sunspot regardless of the presence or absence of a penumbra.
- Move approximately radially outward from sunspots at about 1 km s?1 until they vanish or reach the network.
- Appear with both magnetic polarities from sunspots of single polarities but appear with a net flux of the same sign as the parent sunspot.
- Transport net flux away from the parent sunspots at the same rates as the flux decay of the sunspots.
- Tend to appear in opposite polarity pairs.
- Appear to carry a total flux away from sunspots several times larger than the total flux of the sunspots.
- Produce only a very faint emmission in the core of Hα.
8.
R. K. Zhigalkin G. V. Rudenko N. N. Stepanian V. G. Fainshtein N. I. Shtertser 《Bulletin of the Crimean Astrophysical Observatory》2008,104(1):67-78
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. 相似文献
9.
Markus J. Aschwanden 《Journal of Astrophysics and Astronomy》2008,29(1-2):3-16
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:
- The solar neutrino problem
- Structure of the solar interior (helioseismology)
- The solar magnetic field (dynamo, solar cycle, corona)
- Hydrodynamics of coronal loops
- MHD oscillations and waves (coronal seismology)
- The coronal heating problem
- Self-organized criticality (from nanoflares to giant flares)
- Magnetic reconnection processes
- Particle acceleration processes
- Coronal mass ejections and coronal dimming
10.
Franca Chiuderi Drago 《Solar physics》1970,13(2):357-371
The results of the total solar eclipse of November 12, 1966, observed at 8 different wave-lengths between 3 and 21 cm, are studied and the spectrum of two active regions present on the disk is deduced. It is shown that the observed increase of the flux of the most intense source in the range 3–10 cm is due to geometrical effects. Neglecting the influence of the magnetic field, the following quantities are deduced.
- the mean and central temperature of the coronal condensation.
- the ∫ corona N 2dh (N = electron density).
11.
Franz-Ludwig Deubner 《Solar physics》1971,17(1):6-20
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:
- 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.
- 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.
- Aperiodic horizontal motions are observed in the photospheric layers, which are probably connected with the process of excitation of the oscillatory field.
- 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.
- The correlation obtained by previous observers between vertical stationary motions, the chromospheric network and magnetic fields in particular is confirmed.
12.
N. R. Sheeley Jr. 《Solar physics》1980,66(1):79-87
Unique timelapse sequences of Skylab/ATM spectroheliograms reveal the following characteristics of normal (i.e. non-flare) loop structures in the solar atmosphere:
- At the 0.5 × 106 K temperature of Ne vii, emission is concentrated into individual spiky structures that project 104–105 km from their magnetic footpoints and live on the order of 30 min.
- At the 1.0 × 106 K temperature of Mg ix, the individual spikes are more diffuse, and have greater lengths and longer lifetimes (～ 1.5 hr) than their 0.5 × 106 K counterparts. Perhaps for this reason, more 1.0 × 106 K loops are visible than 0.5 × 106 K loops at any given time.
- At the 2.0 × 106 K temperature of Fe xv, emission is confined to a number of relatively diffuse and irregularly shaped features whose collective patterns define closed field volumes in and between active regions. Although the individual features evolve on a time scale of roughly 6 hr, their collective patterns last for several days or more. Unlike the 0.5 × 106 K features, the 2.0 × 106 K features never form as a linear extension along an apparent magnetic field line, but seem to brighten and fade in place.
13.
Yu. A. Nagovitsyn 《Astrophysical Bulletin》2008,63(1):43-55
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. 相似文献
14.
E. Tandberg-Hanssen P. Kaufmann E. J. Reichmann D. L. Teuber R. L. Moore L. E. Orwig H. Zirin 《Solar physics》1984,90(1):41-62
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:
- The overall magnetic field configuration in which the flare occurred was a fairly simple, closed arch containing nonpotential substructure.
- The flare occurred spontaneously within the arch; it was not triggered by emerging magnetic flux.
- 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.
- The impulsive energy release most likely occurred in the upper part of the arch; it had three immediate products:
- 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.
- 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.
- 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.
15.
I. Lerche 《Astrophysics and Space Science》1977,48(2):335-356
We investigate the ‘equilibrium’ and stability of spherically-symmetric self-similar isothermal blast waves with a continuous post-shock flow velocity expanding into medium whose density varies asr ?ω ahead of the blast wave, and which are powered by a central source (a pulsar) whose power output varies with time ast ω?3. We show that:
- for ω<0, no physically acceptable self-similar solution exists;
- for ω>3, no solution exists since the mass swept up by the blast wave is infinite;
- ? must exceed zero in order that the blast wave expand with time, but ?<2 in order that the central source injects a finite total energy into the blast wave;
- for 3>ωmin(?)>ω>ωmax(?)>0, where $$\begin{gathered} \omega _{\min } (\varphi ){\text{ }} = {\text{ }}2[5{\text{ }} - {\text{ }}\varphi {\text{ }} + {\text{ }}(10{\text{ }} + {\text{ 4}}\varphi {\text{ }} - {\text{ 2}}\varphi ^2 )^{1/2} ]^2 [2{\text{ }} + {\text{ (10 }} + {\text{ 4}}\varphi {\text{ }} - {\text{ 2}}\varphi ^2 {\text{)}}^{{\text{1/2}}} ]^{ - 2} , \hfill \\ \omega _{\max } (\varphi ){\text{ }} = {\text{ }}2[5{\text{ }} - {\text{ }}\varphi {\text{ }} - {\text{ }}(10{\text{ }} + {\text{ 4}}\varphi {\text{ }} - {\text{ 2}}\varphi ^2 )^{1/2} ]^2 [2{\text{ }} - {\text{ (10 }} + {\text{ 4}}\varphi {\text{ }} - {\text{ 2}}\varphi ^2 {\text{)}}^{{\text{1/2}}} ]^{ - 2} , \hfill \\ \end{gathered} $$ two critical points exist in the flow velocity versus position plane. The physically acceptable solution must pass through the origin with zero flow speed and through the blast wave. It must also pass throughboth critical points if \(\varphi > \tfrac{5}{3}\) , while if \(\varphi< \tfrac{5}{3}\) it must by-pass both critical points. It is shown that such a solution exists but a proper connection at the lower critical point (for ?>5/3) (through whichall solutions pass with thesame slope) has not been established;
- for 3>ω>ωmin(?) it is shown that the two critical points of (iv) disappear. However a new pair of critical points form. The physically acceptable solution passing with zero flow velocity through the origin and also passing through the blast wave mustby-pass both of the new critical points. It is shown that the solution does indeed do so;
- for 3>ωmin(?)>ωmax(?)>ω it is shown that the dependence of the self-similar solution on either ω or ? is non-analytic and therefore, inferences drawn from any solutions obtained in ω>ωmax(?) (where the dependence of the solutionis analytic on ω and ?) are not valid when carried over into the domain 3>ωmin(?)>ωmax(?)>ω;
- all of the physically acceptable self-similar solutions obtained in 3>ω>0 are unstable to short wavelength, small amplitude but nonself-similar radial velocity perturbations near the origin, with a growth which is a power law in time;
- the physical self-similar solutions are globally unstable in a fully nonlinear sense to radial time-dependent flow patterns. In the limit of long times, the nonlinear growth is a power law in time for 5<ω+2?, logarithmic in time for 5>ω+2?, and the square of the logarithm in time for 5=ω+2?.
16.
N. Hiotelis 《Astrophysics and Space Science》2008,315(1-4):191-200
We used merger trees realizations, predicted by the extended Press-Schechter theory, in order to study the growth of angular momentum of dark matter haloes. Our results showed that:
- The spin parameter λ′ resulting from the above method, is an increasing function of the present day mass of the halo. The mean value of λ′ varies from 0.0343 to 0.0484 for haloes with present day masses in the range of 109h?1 M ⊙ to 1014h?1 M ⊙.
- The distribution of λ′ is close to a log-normal, but, as it is already found in the results of N-body simulations, the match is not satisfactory at the tails of the distribution. A new analytical formula that approximates the results much more satisfactorily is presented.
- The distribution of the values of λ′ depends only weakly on the redshift.
- The spin parameter of an halo depends on the number of recent major mergers. Specifically the spin parameter is an increasing function of this number.
17.
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:
- T(r) ; T0 (isothermal case)
- T(r) ; T0(r/r0)1/2
- T(r) ; T0(r/r0)-1
- T(r) ; T0(r/r0)-2
- T(r) ; T0(r/r0)-3
18.
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:
- 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°.
- 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.
- The maximum number of polar faculae was reached at the activity minimum (1975–1976).
- The last polar faculae were observed in the second half of 1978 at the latitudes from 70° to 80°.
19.
Joseph V. Hollweg 《Solar physics》1978,56(2):305-333
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:
- 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.
- 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.
- ¦δv¦ and ¦δB¦ also display resonant peaks.
- 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.
- The dependences of ¦δv¦ and ¦δB¦ on height are reduced by finite wavelength effects, except near the wave source where they are enhanced.
- 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.
- Below the low corona most wave energy is kinetic, except near the base where it becomes mostly magnetic at the resonances.
- ?0 < δv 2 > v A or < δB 2 > v A/4π are not good estimators of the energy flux.
- 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.
20.
McWhirter et al. (1975) have presented a standard model for the transition region and inner corona that matches with the Harvard Smithsonian Reference Atmosphere. They assume an open field line configuration and solve numerically the equations of energy and hydrostatic equilibrum. The purpose of the present paper is to generalise their model for the temperature and density as functions of height in several ways and, in particular, to determine the temperature maxium and its location. The effect of varying the following characteristics of the model is determined:
- Boundary conditions on temperature and density;
- magnitude of the heating;
- form of the heating term;
- divergence of the field lines;
- presence of subsonic flows, either upward or downward.