The structure of the solar flare stream magnetic field

The structure of the interplanetary magnetic field within the flare streams as well as associated variations of the geomagnetic disturbancy are considered. It is shown that in the main body of the flare stream the magnetic field is determined by the configuration of the large scale magnetic field on the Sun at the flare region. Within the head part of the flare stream the magnetic field represents by itself the compressed field of the background solar wind and hence is determined by the distribution of the super large scale solar magnetic field outside the flare region.A certain asymmetry in the parameters of the magnetic field within the streams associated with geoeffective and non-effective flares is shown to exist.     

Fluxtube magnetic field measurements in a solar flare in comparison with non-flare regions

Measurements of the fluxtube field strength B and filling α in solar flares, active regions and faculae have been analyzed. To estimate the values of B, Stokes V peak separations of the Fe I 5247 Å and Fe I 5250 Å lines have been used. It was found that the value of B in an In-flare (˜ 1.1 kG) was slightly smaller than that in faculae (˜ 1.3 kG) and a non-flare active region (˜ 1.4kG). On the other hand, in a more power 2b-flare the value of B was in the range from. 1.1 kG (the start of the flare) to 1.55kG. (during the peak). Thus the values of the field strength of flares somewhat differ from those both of faculae and active regions. The magnetic filling factor is approximately equal in flares (0.2-0.40) and active regions (0.3) but several times larger as compared with faculae.     

Small-scale magnetic flux tube diagnostics in a solar flare

I ±V profiles of the Fei 5247 and 5250 lines in the 2B flare of June 16, 1989 have been analyzed. A bright knot of the flare outside the sunspot where the central intensity of H reached a peak value of 1.4 (relative to the continuum) has been explored. The Fei 5250/Fei 5247 magnetic line ratio based on the StokesV peak separations of these lines at five evolutionary phases of the flare (including the start of the flare, the flash phase, the peak and 16 min after the peak) has been analyzed. It was found that the StokesV peak separation for the Fei 5250 line was systematically larger than that of the Fei 5247 line. This is evidence for the presence in the flare of small-scale flux tubes with kG fields. The flux tube magnetic field strength was about 1.1 kG at the start of the flare and during the flash phase, 1.55 kG during the peak, and 1.38 kG 16 min after the peak. The filling factor,, appears to decrease monotonically during the flare.     

Observations of fine time structure in solar flare hard X-ray bursts

Hard X-ray (?100 keV) time histories of solar flares which occurred on 1978 December 4 and 1979 February 18 are presented. The first flare was observed by 3 identical instruments from near-earth orbit (Prognoz 7) and interplanetary space (Venera 11 and 12). Fine time structure is present down to the 55 ms level for the e-folding rise and fall times. These data may be used to localize the emission region by the method of arrival time analysis.     

Coronal general magnetic field evolution as a new parameter of the solar cycle

The magnetic field lines of the corona associated with the solar-cycle surface general magnetic field are calculated by a potential-field approximation to study the solar-cycle evolution of the geometry of the coronal field. The surface field evolution used here is the radial field evolution, predicted by a model of the solar cycle driven by the dynamo action of the global convection, and justified observationally using Mount Wilson magnetic synoptic chart data. The evolution of the calculated coronal general field is now good for comparison with observations and shows the following. (i) The field of the polar and high-altitude corona has dipolar structure in almost all phases of the solar cycle except in a short time interval around maximum phase despite the quadrupolar structure of the general magnetic field at the surface; quadrupolar field forms loop-like structure in the lower corona. The almost-dipolar structure of the polar and high-latitude corona and the loop formation of the equatorial lower corona explain the appearance of the undisturbed minimum corona observed at eclipses. (ii) The polar field lines are directed almost radially at the minimum phase, which should be responsible for polar plumes. The field lines slowly open up to participate in the loop-like structure of the equatorial lower corona, and rapidly change their structure and polarity at the maximum phase, to resume the almost radial configuration slowly, (iii) During the rapidly changing maximum phase, the field lines do not penetrate deep into the interplanetary space resulting in the absence of polar plumes and the appearance of the circular corona- the maximum corona. In this phase, the coronal field should not be approximated by a dipole field. The surface field evolution which can explain such behaviors of the corona is characteristic of the solar-cycle process dominated by the latitudinal gradient of the differential rotation. If the radial gradient dominated in the subsurface process, the coronal evolution would look quite different and would show latitudinal propagation of enhancement of activity. Although nonaxisymmetric features should be superposed on the axisymmetric general field to express the real corona, the general field can be a basic coronal field in studying long-term interaction between the convection zone and the interstellar space especially in studying the magnetic braking of the solar rotation.     

A magnetic bald-patch flare in solar active region 11117

With SDO observations and a data-constrained magnetohydrodynamics(MHD)model,we identify a confined multi-ribbon flare that occurred on 2010 October 25 in solar active region 11117 as a magnetic bald patch(BP)flare with strong evidence.From the photospheric magnetic field observed by SDO/HMI,we find there are indeed magnetic BPs on the polarity inversion lines(PILs)which match parts of the flare ribbons.From the 3D coronal magnetic field derived from an MHD relaxation model constrained by the vector magnetograms,we find strikingly good agreement of the BP separatrix surface(BPSS)footpoints with the flare ribbons,and the BPSS itself with the hot flaring loop system.Moreover,the triggering of the BP flare can be attributed to a small flux emergence under the lobe of the BPSS,and the relevant change of coronal magnetic field through the flare is reproduced well by the pre-flare and post-flare MHD solutions,which match the corresponding pre-and post-flare AIA observations,respectively.Our work contributes to the study of non-typical flares that constitute the majority of solar flares but which cannot be explained by the standard flare model.     

Microwave burst spectra and solar flare magnetic fields

Microwave burst spectra are compared with the position, within the active region, of their associated flares observed in H. The magnetic fields predicted by Takakura's burst model (1972) are found to be in reasonable agreement with the fields expected at the flare locations.Operated by the Association of Universities for Research in Astronomy, Inc., under contract with the National Science Foundation.     

Observations of small-scale solar magnetic fields

Spectrograms, obtained during moments of good seeing with the high spatial resolution afforded by the 80-cm solar image at the Kitt Peak National Observatory, show the following:

1. (1)
   Magnetic fields of several hundred gauss occur in tiny areas easily as small as 500 km in extent in regions of the solar surface sometimes well removed from areas of sunspot activity.     
   
   

Variations of magnetic fields and electric currents associated with a solar flare

The high-resolution vector magnetograms obtained with the solar telescope magnetograph of the Beijing Astronomical Observatory of the active region AR 4862 on 7 October, 1987, close before and after a solar flare, were used to calculate the electric current densities in the region. Then the relations between the flare and the magnetic fields as well as the electric currents were studied. The results are: (i) the transverse magnetic fields, and hence the longitudinal electric currents in the region before and after the flare, are evidently different, while the longitudinal magnetic fields remain unchanged; (ii) this confirms the result obtained previously that the flare kernels coincide with the peaks of longitudinal electric density in active regions; (iii) the close relation between the flare kernels and the electric currents indicates that the variations of the transverse magnetic fields and the longitudinal electric currents arise not from the general global evolution of the active region, but from the flare. These results tend to the conclusion that the triggering of a solar flare might be related with the plasma instability caused by the surplus longitudinal electric currents at some local regions in the solar atmosphere.     

Spatial and temporal evolution of soft and hard X-ray emission in a solar flare

We study the spatial and temporal characteristics of the 3.5 to 30.0 keV emission in a solar flare on April 10, 1980. The data were obtained by the Hard X-ray Imaging Spectrometer aboard the Solar Maximum Mission Satellite. It is complemented in our analysis with data from other instruments on the same spacecraft, in particular that of the Hard X-ray Burst Spectrometer.Key results of our investigation are: (a) Continuous energy release is needed to substain the increase of the emission through the rising phase of the flare, before and after the impulsive phase in hard X-rays. The energy release is characterized by the production of hot (5 × 10⁷ T 1.5 × 10⁸ K) thermal regions within the flare loop structures. (b) The observational parameters characterizing the impulsive burst show that it is most likely associated with non-thermal processes (particle acceleration). (c) The continuous energy release is associated with strong chromospheric evaporation, as evidenced in the spectral line behavior determined from the Bent Crystal Spectrometer data. Both processes seem to stop just before flare maximum, and the subsequent evolution is most likely governed by the radiative cooling of the flare plasma.     

Predicting the evolution of photospheric magnetic field in solar active regions using deep learning

The continuous observation of the magnetic field by the Solar Dynamics Observatory(SDO)/Helioseismic and Magnetic Imager(HMI) produces numerous image sequences in time and space.These sequences provide data support for predicting the evolution of photospheric magnetic field. Based on the spatiotemporal long short-term memory(LSTM) network, we use the preprocessed data of photospheric magnetic field in active regions to build a prediction model for magnetic field evolution. Because of the elaborate learning and memory mechanism, the trained model can characterize the inherent relationships contained in spatiotemporal features. The testing results of the prediction model indicate that(1) the prediction pattern learned by the model can be applied to predict the evolution of new magnetic field in the next 6 hours that have not been trained, and predicted results are roughly consistent with real observed magnetic field evolution in terms of large-scale structure and movement speed;(2) the performance of the model is related to the prediction time; the shorter the prediction time, the higher the accuracy of the predicted results;(3) the performance of the model is stable not only for active regions in the north and south but also for data in positive and negative regions. Detailed experimental results and discussions on magnetic flux emergence and magnetic neutral lines finally show that the proposed model could effectively predict the large-scale and short-term evolution of the photospheric magnetic field in active regions. Moreover, our study may provide a reference for the spatiotemporal prediction of other solar activities.     

The correlation of solar flare production with magnetic energy in active regions

An investigation of 531 active regions was made to determine the correlation between energy released by flares and the available energy in magnetic fields of the regions. Regions with magnetic flux greater than 10²¹ maxwell during the years 1967–1969, which included sunspot maximum, were selected for the investigation. A linear regression analysis of flare production on magnetic flux showed that the flare energy is correlated with magnetic energy with a coeificient of correlation of 0.78. Magnetic classification and field configuration also significantly affect the production of flares.This work was supported by the Aerospace Sponsored Research Program.     

Strong magnetic field in W75N OH maser flare

A flare of OH maser emission was discovered in W75N in 2000. Its location was determined with the Very Long Baseline Array (VLBA) to be within 110 au from one of the ultracompact H  ii regions, Very Large Array 2 (VLA2). The flare consisted of several maser spots. Four of the spots were found to form Zeeman pairs, all of them with a magnetic field strength of about 40 mG. This is the highest ever magnetic field strength found in OH masers, an order of magnitude higher than in typical OH masers. Three possible sources for the enhanced magnetic field are discussed: (i) the magnetic field of the exciting star dragged out by the stellar wind; (ii) the general interstellar field in the gas compressed by the magnetohydrodynamic shock; and (iii) the magnetic field of planets which orbit the exciting star and produce maser emission in gaseous envelopes.     

The large-scale solar magnetic field

The large-scale photospheric magnetic field, measured by the Mt. Wilson magnetograph, has been analyzed in terms of surface harmonics (P _n ^m )()cosm and P _n ^m ()sinm) for the years 1959 through 1972. Our results are as follows. The single harmonic which most often characterized the general solar magnetic field throughout the period of observation corresponds to a dipole lying in the plane of the equator (2 sectors, n = m = 1). This 2-sector harmonic was particularly dominant during the active years of solar cycles 19 and 20. The north-south dipole harmonic (n = 1, m = 0) was prominent only during quiet years and was relatively insignificant during the active years. (The derived north-south dipole includes magnetic fields from the entire solar surface and does not necessarily correlate with either the dipole-like appearance of the polar regions of the Sun or with the weak polar magnetic fields.) The 4-sector structure (n = m = 2) was prominent, and often dominant, at various times throughout the cycle. A 6-sector structure (n = m = 3) occasionally became dominant for very brief periods during the active years. Contributions to the general solar magnetic field from harmonics of principal index 4 n 9 were generally relatively small throughout this entire solar cycle with one outstanding exception. For a period of several months prior to the large August 1972 flares, the global photospheric field was dominated by an n = 5 harmonic; this harmonic returned to a low value shortly after the August 1972 flare events. Rapid changes in the global harmonics, in particular, relative and absolute changes in the contributions of harmonics of different principal index n to the global field, imply that the global solar field is not very deep or that very strong fluid flows connect the photosphere with deeper layers.The National Center for Atmospheric Research is sponsored by the National Science Foundation.     

Enigmas in measurements of solar magnetic field

The mean magnetic field (MMF) of the photosphere of the Sun as a star was measured in 2001?C2010 at the Crimean Astrophysical Observatory using two Fe I absorption lines with ?? = 524.7 nm and ?? = 525.0 nm. The regression coefficient b for 1054 pairs of daily values measured simultaneously on both lines equals 0.82 (a correlation coefficient is 0.94; magnetic field strengths determined by the line with ?? = 525.0 nm are lower than those for the line with ?? = 524.7 nm). However, the b value varied significantly along with phases of the 11-year cycle from 0.88 in 2003 to 0.49 in 2009. It is difficult to ascribe these variations to purely instrumental or solar causes. Moreover, the semiannual value of b decreased with the decrease in the absolute strength of the MMF, which contradicts the model of thin magnetic flux ropes of the photosphere. Similar behavior of b was also observed in the comparison of MMF measured at the Crimean Astrophysical Observatory and Stanford by the line with ?? = 525.0 nm. The inconsistency of the results obtained by these two iron lines on different instruments has been noted. It has been concluded that the variance in and odd behavior of b are predetermined not only by the instrument and the Sun (by the so-called fine structure of the photosphere field), but also by the act of measuring. When recording solar (and stellar) magnetic fields and modeling atmospheric processes, quantum effects have to be taken into account, such as nonlocality, indistinguishability, and the entanglement of photons, as well as that a photon only acquires its properties at the exact moment of its detection. The best approximation to reality can be achieved by averaging the MMF measurements carried out with different magnetographs and in different spectral lines.     

Observations of energetic X-rays and solar cosmic rays associated with the 23 May 1967 solar flare event

On 23 May 1967 energetic (10–50 keV) solar flare X-rays were observed by the OGO-III ion chamber during the period 1808–2100 UT. The time-intensity profile for the X-ray event showed three distinct peaks at 1810, 1841 and 1942 UT. The second peak, which is equivalent to 2.9 × 10^–3 ergs cm^–2sec^–1 above 20 keV, is the largest X-ray burst observed so far by the OGO-I and OGO-III ion chambers. The soft (2–12 Å) X-ray observations reported by Van Allen (1968) also show similar peaks, roughly proportional in magnitude to the energetic X-ray peaks. However, the intensity of energetic X-rays peaked in each case 5–10 min earlier than the soft X-ray intensity indicating a relatively hard photon energy spectrum near the peak of the energetic X-ray emission. The corresponding time-intensity profile for the solar radio emission also showed three peaks in the microwave region nearly coincident with the energetic X-ray peaks. The third radio peak was relatively rich in the metric emission. Beyond this peak both the energetic X-rays and the microwave emission decayed with a time constant of 8 min while the corresponding time constant for the soft X-rays was 43 min. In view of the earlier findings about the energetic X-rays it is indicated that the 23 May solar X-ray event was similar to those observed earlier. During the 23 May event the integral energy flux spectrum at the time of peak intensity is found to be consistent with the form e ^–E/E ₀, E ₀ being about 3.4 and 3.7 keV for the peaks at 1841 and 1942 UT, respectively. Assumption of a similar spectrum during the decay phase indicates that the spectral index E ₀ decreased nearly exponentially with time.The OGO-III ion chamber, which is also sensitive to protons 12 MeV, observed a solar particle event starting at 2100 UT on 23 May. It could not be determined uniquely which of the two principal X-ray peaks was associated with the particle event, and in fact both may have contributed. The particle intensity reached its maximum value at 1003 UT on 25 May 1967. The equivalent peak radiation dosage was 24 R/hour behind the 0.22 g cm^–2 thick aluminum wall of the chamber. This peak radiation dosage was considerably smaller than the maximum dosage (60 R/hour) during the 2 September 1966 solar particle event, the largest event observed so far by the OGO-I and OGO-III satellites. The temporal relationship between the solar X-ray and particle events on 23 May 1967 was similar to that observed in the solar flare events on 7 July 1966, 28 August 1966 and 27 February 1967.     

A kinematic model of a solar flare

Hyder advocated the idea that the optical (H) flares can be identified with the response of the solar chromosphere to an infalling material stream resulting from the disparition brusque of a prominence. Since some flares are observed without any apparent association with infalling streams, in this paper we examine the possibility of identifying the optical flare with the response of the chromosphere to a supersonic disturbance, i.e., a shock, propagating downward. The undisturbed chromosphere is represented by the Harvard-Smithsonian Reference Atmosphere and the evolution of the shock is evaluated with the use of the CCW (Chisnell, Chester, Whitham) approximation based on the theory of characteristics. It is shown that the chromosphere is heated by the shock and that radiation is enhanced, and that the enhanced radiation terminates the shock around the height of the temperature minimum. Numerical results obtained and possible future improvements of this type of study are discussed.The National Center for Atmospheric Research is sponsored by the National Science Foundation.     

Measurement and analysis of magnetic field variation during a class 2b flare

We report a digital analysis of high-time-resolution videomagnetograms taken during a class 2b flare that occurred at 60° east. The data were obtained at the Big Bear Observatory and calibrated by a Mt. Wilson magnetogram. Changes of weak magnetic fields (less than 100 G) with an amplitude from 30 to 100% have been detected over 55% of the optical flare region, apparently taking place at the initial phase of the flare. Statistical considerations suggest a real flare association with most of these changes.H observations show that large changes took place over the footpoints of heavily inclined structures like penumbral fibrils, while smaller changes took place over the plage region. An apparent polarity reversal was found at the feet of erupted fibrils.Based on force-free field calculations these changes can be reasonably explained as a transformation of the current-carrying fields to potential fields which produced large changes in the field line inclination and rotation.Visiting Associate, Summer 1977.