Investigation of five white light flares

On the basis of original observations, five white-light flares (WLF) are investigated. Evidence is given that their emission is located in two points brightening on either side of the lineH _∥ = 0 and lying at the foot points of chromospheric loops. The area of WLF is ≈ 5 × 10^?6 hemisphere, i.e. ≈ 0.007 of Hα flare area; the intensity of WLF is sometimes twice that of the background at the center of the disk. WLF are resolved into more bright and fine knots of ≈ 2′ in diameter. The position of WLF coincides with the brightest knots of Hα flares which are characterized by wide wings with rapid increase and decrease. According to our estimates, the full output of the energy of a flare in the continuum and Hα are comparable; but, the energy emitted in integral light in time-unit through area-unit is by 2 orders of magnitude larger than the energy in monochromatic light.     

The June 15, 1991 and June 26, 1999 white-light flares

Observational properties of two white-light flares (WLFs), on June 15, 1991, and June 26, 1999, are presented and compared. This is of particular interest, because the former was one of the most intense flares of X-ray class X12, while the latter was a compact flare of class M2.3. Significant differences between some flare parameters (GOES class, H_α classification, the number of WLF kernels and their location in the sunspot group, the size and duration of the WLF emission, and the peak flux density of the microwave emission) have been found. However, both these events had approximately the same powers of the emission per unit area in continuum near 658.0 nm: E = 1.5 × 10⁷ and 1.1. × 10⁷ erg cm^?2 s^?1 nm^?1. There is generally a good temporal coincidence between the microwave and hard X-ray emissions and the WLF emission during the impulsive phase, but the light curve of the WLF emission on June 26, 1999, shows a stronger correlation with the X-ray emission in the energy range 14–23 keV. Both flares can be classified by their spectral characteristics as type I white-light flares.     

Electron precipitation and mass motion in the 1991 June 9 white-light flare

We use H line profiles as a diagnostic of mass motion and nonthermal electron precipitation in the white-light flare (WLF) of 1991 June 9 01:34 UT. We find only weak downflow velocities (10 km s^–1) at the site of white-light emission, and comparable velocities elsewhere.We also find that electron precipitation is strongest at the WLF site. We conclude that continuum emission in this flare was probably caused by nonthermal electrons and not by dynamical energy transport via a chromospheric condensation.     

The occurrence frequency of white-light flares

We derive an occurrence frequency for white-light flares (WLF) of 15.5 ± 4.5 yr^?1 during a 2.6 year period following the maximum of solar cycle 21. This compares with a frequency 5–6 yr^?1 derived by McIntosh and Donnelly (1972) during solar cycle 20. We find that the higher frequency of the more recently observed WLFs is due to the availability of patrol data at shorter wavelengths (λ ? 4000 Å), where the contrast of the flare emission is increased; the improved contrast has allowed less energetic (and hence more frequently occurring) events to be classified as WLFs. We find that sufficient conditions for the occurrence of a WLF are: active region magnetic class = delta; sunspot penumbra class = K, with spot group area ≥ 500 millionths of the solar hemisphere; 1–8 Å X-ray burst class ≥ X2.     

Observational characteristics of the white-light flare of August 9, 2011

We analyzed the monochromatic H_α and spectral (within a range of 6549–6579 Å) observational data for the 2B/X6.9 flare of August 9, 2011, that produced emission in the optical continuum. The morphology and evolution of the H_α flare and the position, time evolution, spectrum, and energetics of the white-light flare (WLF) kernels were studied. The following results were obtained: the flare erupted in the region of collision of a new and rapidly growing and propagating magnetic flux and a preexisting one. This collision led to a merger of two active bipolar regions. The white-light flare had a complex structure: no less than five kernels of continuum emission were detected prior to and in the course of the impulsive flare phase. Preimpulsive and impulsive white-light emission kernels belonged to different types (types II and I, respectively) of white-light flares. A close temporal agreement between the white-light emission maxima and the microwave emission peak was observed for the impulsive white-light emission kernels. The maximum flux, luminosity, and total energy emitted by the brightest impulsive WLF kernel equaled 1.4 × 10¹⁰ ergs cm^?2 s^?1, 1.5 × 10²⁷ ergs/s, and 5 × 10²⁹ ergs, respectively. The H_α profiles within the impulsive WLF kernels had broad wings (with a total extent of up to 26 Å and a half-width of up to 9 Å) and self-reversed cores. The profiles were symmetrical, but were shifted towards the red side of the spectrum. This is indicative of a downward motion of the entire emitting volume with a radial velocity of several tens of km/s. The intensity pattern in the wings did not correspond to the Stark one. The profiles were broadened by nonthermal turbulent motions with velocities of 150–300 km/s. The observed H_α profiles were analyzed and compared in their features to the profiles calculated for an intense heating of the chromosphere by nonthermal electrons accompanied by the development of a chromospheric condensation propagating downward. We came to the conclusion that the analyzed flare exhibited spectral features that may not be readily explained within the framework of chromosphere heating by a beam of nonthermal electrons.     

Ni I 6768 Å line profile variations during a solar flare and their effect on the SOHO/MDI magnetic field measurements

Observational data on the Ni I 6768 Å line profile variations during the impulsive and post-impulsive phases of the July 18, 2002 while light flare (WLF) in the kernel of WLF emission and in other flare kernels are presented. The line profiles at the sites of intense photospheric motions in active regions are also studied. The effect of the observed Ni I 6768 Å line profile variations on the SOHO/MDI magnetic field measurements is estimated. The following conclusions have been reached. (1) The thermodynamic structure of the photo-spheric layers changes significantly during the flare. As a result, the Ni I line profile changes, particularly at the site of WLF emission. At this time, the line depth decreases significantly, but the line does not show any emission reversal. Subsequently, a relatively slow return to the conditions of an undisturbed photosphere is observed. (2) The technique of SOHO/MDI magnetic field measurements is insensitive to such line variations. Therefore, the detected variations during the flare did not result in any noticeable errors in the MDI longitudinal magnetic field measurements. (3) The line profile is broadened, shifted as a whole, and asymmetric at the sites of active regions where intense photospheric motions appear. In the MDI measurements, such changes in the profile lead to an underestimation of the magnetic field by approximately 10% if the line-of-sight velocity of the photo-spheric ejection is about 1.6 km s^?1.     

Force-free magnetic fields may have a "loss of equilibrium"

By comparing the light curves in optical, hard x-ray, and soft x-ray wavelengths for 8 well-observed flares, we confirm previous results indicating that the white light flare (WLF) is associated with the flare impulsive phase. The WLF emission peaks within secondsafter the associated hard x-ray peak, and nearly two minutesbefore the 1–8 soft x-ray peak. It is further shown that the peak power in nonthermal electrons above 50 keV is typically an order of magnitudelarger, and the power in 1–8 soft x-rays radiated over 2 strdn at the time of the WLF peak is an order of magnitudesmaller, than the peak WLF power.Operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation. Partial support for the National Solar Observation is provided by the USAF under a Memorandum of Undestanding with the NSF.     

High resolution observations of white‐light emissions from the opacity minimum during an X‐class flare

Using high cadence, high resolution near infrared (NIR) observations of the X10 white‐light flare (WLF) on 2003 October 29, we investigated the evolution of the core‐halo structure of white‐light emission during the two‐second period flare peak. We found that size and intensity of the halo remained almost constant in the range of 10 Mm². However, the core area was very compact and expanded rapidly from about 1 Mm² to 4 Mm². At the same time, the total emission of the core increased nearly twenty times. This distinct behavior indicates that different heating mechanisms might be responsible for core and halo emissions. In addition to the temporal analysis, we compared the intensity enhancements of the flare core and halo. The result shows that the halo contrast increased by about 8% compared to the flare‐quiet region, which could be explained by a combination of direct‐heating and backwarming models (© 2010 WILEY‐VCH Verlag GmbH & Co. KGaA, Weinheim)     

A Yohkoh search for “black-light flares”

Calculations which predict that a phenomenon analogous to stellar negative pre-flares could also exist on the Sun were published by Hénouxet al. (1990), and Aboudarhamet al., (1990), who showed that at the beginning of a solar white-light flare (WLF) event an electron beam can cause a transient darkening before the WLF emission starts, under certain conditions. They named this event a black light flare (BLF). Such a BLF event should appear as diffuse dark patches lasting for about 20 seconds preceding the WLF emission, which would coincide with intense and impulsive hard X-ray bursts. The BLF location would be at (or in the vicinity of) the forthcoming bright patches. Their predicted contrast depends on the position of the flare on the solar disc and on the wavelength band of the observation.TheYohkoh satellite provided white-light data from the aspect camera of the SXT instrument (Tsunetaet al., 1991), at 431 nm and with a typical image interval of 10–12 s. We have studied nine white-light flares observed with this instrument, with X-ray class larger than M6. We have found a few interesting episodes, but no unambiguous example of the predicted BLF event. This study, although the best survey to date, was not ideal from the observational point of view. We therefore encourage further searches. Successful observations of this phenomenon on the Sun would greatly strengthen our knowledge of the lower solar atmosphere and its effects on solar luminosity variations.     

The hydrogen emission spectrum in three white-light flares

We present spectral data for three white-light flares (WLFs) showing Balmer continuum at wavelengths 3700 Å. These flares also have a weaker continuum extending toward longer wavelengths, from which, in one flare where this continuum is sufficiently bright, we are able to identify a Paschen jump near 8500 Å. The presence of the latter suggests that the Paschen continuum may be a substantial contributor to the WLF continuum at visible wavelengths. We note the possibility, therefore, that the entire continuum of this particular flare may be dominated by H _fb emission.In all three flares the head of the Balmer continuum, as well as the head of the Paschen continuum in the flare where it was identified, is advanced toward longer wavelengths as a result of the blending of the hydrogen emission lines of the respective series. The principal quantum number of the last resolvable line of the Balmer or Paschen series is approximately 16. The electron density, as measured from the halfwidths of the high Balmer lines in two of the flares, is approximately 5 × 10¹³ cm^–3. Due to possible misplacements of the spectrograph slit, however, the electron density in the brightest kernels of the WLFs may not have been obtained.Operated by the Association of Universities for Research in Astronomy, Inc. under contract AST 78-17292 with the National Science Foundation.     

Occurrence of a White-Light Flare Over a Sunspot and the Associated Change of Magnetic Flux

A white-light flare (WLF) on 10 March 2001 was well observed in the Hα line and the Ca ii λ8542 line using the imaging spectrograph installed on the Solar Tower Telescope of Nanjing University. Three small sunspots appeared in the infrared continuum image. In one sunspot, the infrared continuum is enhanced by 4–6% compared to the preflare value, making the sunspot almost disappear in the continuum image for about 3 min. A hard X-ray (HXR) source appeared near the sunspot, the flux of which showed a good time correlation with the profile of the continuum emission. In the sunspot region, both positive and negative magnetic flux suffered a substantial change. We propose that electron precipitation followed by radiative back-warming may play the chief role in heating the sunspot. The temperature rise in the lower atmosphere and the corresponding energy requirement are estimated. The results show that the energy released in a typical WLF is sufficient to power the sunspot heating.     

太阳射电IV型爆发及伴随现象与白光耀斑的关系

通过1991年6月6日共生太阳白光耀斑（WLF）的射电运动IV型爆发及其伴随现象（包括耀斑后环、爆发衰减相的射电脉动、多波段射电辐射和太阳物质抛射等）观测资料的分析,定性地探讨了WLF的起源、加热机制和发射地点的问题．假设了WLF和射电运动IV型射电爆发可能有共同起源的低日冕电子加速区,讨论了WLF的能量传输可能是通过二步加速过程,即来自低日冕的非热电子沉降能量于色球层,产生色球层的压缩波或向下的辐射场进而使上光球层温度增加导致WLF此外,提出WLF可能会伴有耀斑后环和射电精细结构的对应物．     

X-Ray Emission from Comet McNaught-Hartley (C/1999 T1)

Comet McNaught-Hartley was observed in five 1-h exposures on January 8-14 2001 using the advanced CCD imaging spectrometer on board the Chandra X-ray Observatory. The X-ray image of the comet does not show a crescent-like shape. The brightest region is offset from the nucleus between the sunward and comet velocity directions. The comet mean X-ray luminosity is equal to 7.8×10¹⁵ erg s⁻¹ for photon energy E>150 eV and aperture ρ=1.5×10⁵ km where the comet X-ray brightness exceeds 20% of the peak value. Gas production rate was 10²⁹ s⁻¹ during the observations, and the efficiency of X-ray excitation was equal to 4×10⁻¹⁴ erg AU^3/2. Day-to-day variations in X-rays reached a factor of 5. The strongest short-term variation was by a factor of 1.75 for 1600 s. This variation may be explained by a decline in the solar-wind flux by the same factor in ≈800 s. The comet and Earth were seeing different faces of the Sun, and time delay in the solar-wind events on the Earth and the comet was long, equal to 6 days. The best correlation between the comet X-ray luminosity and the solar-wind proton density is for the time delay of 5.5 days and may be explained by the higher velocity of heavy ions.Careful background subtraction made it possible to extract the comet spectrum from 150 to 1000 eV. No signal was detected at E>1000 eV, and a 3σ upper limit to any emission with E>1000 eV is 0.3% of the photon emission at 150-1000 eV. The best χ²-fit model to the spectrum consists of nine narrow emission features. The emission energies and intensities are in good agreement with a charge exchange spectrum calculated by us for the slow solar wind. Using this spectrum, we identify the observed emissions as (Ne⁷⁺+Mg⁷⁺+Mg⁸⁺) at 195 eV, (Mg⁸⁺+Mg⁹⁺+Si⁸⁺) at 250 eV, C⁵⁺ at 370 and 460 eV, O⁶⁺ at 560 eV, O⁷⁺ at 650, 780, and 840 eV, and Ne⁸⁺ at 940 eV. X-ray spectroscopy of comets may be used to diagnose the solar-wind composition and its interaction with comets.     

Observations of the gamma-ray flux from the galaxy Mk 501

We present two-year-long observations of the flux of very-high-energy (～10¹² eV) gamma rays from the active galactic nucleus Mk 501 performed with a Cherenkov detector at the Crimean Astrophysical Observatory. A gamma-ray flux from the object was shown to exist at confidence levels of 11 and 7 standard deviations for 1997 and 1998, respectively. The flux varied over a wide range. The mean flux at energies >10¹² eV, as inferred from the 1997 and 1998 data, is (5.0±0.6)×10^?11 and (3.7±0.6)×10^?11 cm^?2 s^?1, respectively. The errors are the sum of statistical observational and modeling errors. The mean power released in the form of gamma rays is ～2×10⁴³ erg s^?1 sr^?1.     

Airborne interferometric measurement of the infrared spectrum of Jupiter from 6 to 14 μm

A new spectrum of Jupiter from 700 to 1600 cm^?1 was obtained with an interferometric experiment using the 91.5 cm telescope of the NASA Airborne Infrared Observatory. The spectral resolution is 10 cm^?1 and the signal-to-noise ratio is 30 at 900 cm^?1. NH₃ absorption lines are observed between 820 and 1020 cm^?1. The 1306 cm^?1ν₄CH₄ band strongly appears in emission at a temperature of at least 145° K. The Jovian brightness temperature between 1400 and 1600 cm^?1, according to our measurement, is lower than 170° K.     

Observations of the gamma-ray flux from the X-ray source Cyg X-3 in 1994–1995

We present the observations of Cygnus X-3 carried out with the GT-48 gamma-ray telescope at the Crimean Astrophysical Observatory in 1994–1995. The mean gamma-ray flux at energy E>10¹² eV is shown to be approximately equal to 1.3×10^?11 cm^?2 s^?1. The flux in 1994 was much lower than that in 1995, being (6.2±2.6)×10^?12 cm^?2 s^?1; i.e., it was statistically insignificant. The flux in 1995 was (2.7±0.7)×10^?11 cm^?2 s^?1. Thus, the very high energy gamma-ray emission from Cyg X-3 is variable. These measurement results can be used to obtain upper limits on the flux from Cyg X-3 in 1994–1995.     

Structure of radio sources at wavelengths 10.5 and 12.0 cm from observations of the Solar eclipse on March 29, 2006

The eclipse observations were performed at the Laboratory of Radio Astronomy of the CrAO in Katsiveli with stationary instrumentation of the Solar Patrol at wavelengths of 10.5 and 12.0 cm. The data obtained were used to determine the brightness temperature of the undisturbed Sun at solar activity minimum between 11-year cycles 23 and 24: T _d10.5 = (43.7 ± 0.5) × 10³ K at 10.5 cm and T _d12.0 = (51.8 ± 0.5) × 10³ K at 12.0 cm. The radio brightness distribution above the limb group of sunspots NOAA 0866 was calculated. It shows that at both wavelengths the source consisted of a compact bright nucleus about 50 × 10³ km in size with temperatures T _b10.5 = 0.94 × 10⁶ K and T _b12.0 = 2.15 × 10⁶ K located, respectively, at heights h _10.5 = 33.5 × 10³ km and h _12.0 = 43.3 × 10³ km above the sunspot and an extended halo with a temperature T _b = (230–300) × 10³ K stretching to a height of 157 × 10³ km above the photosphere. The revealed spatial structure of the local source is consistent with the universally accepted assumption that the radiation from the bright part of the source is generated by electrons in the sunspot magnetic fields at the second-third cyclotron frequency harmonics and that the halo is the bremsstrahlung of thermal electrons in the coronal condensation forming an active region. According to the eclipse results, the electron density near the upper boundary of the condensation was N _e ≈ 2.3 × 10⁸ cm^?3, while the optical depth was τ ≈ 0.1 at an electron temperature T _e ≈ 10⁶ K. Thus, the observations of the March 29, 2006 eclipse have allowed the height of the coronal condensation at solar activity minimum to be experimentally determined and the physical parameters of the plasma near its upper boundary to be estimated.     

Calculated photoelectron pitch angle and energy spectra

Calculations of the steady-state photoelectron energy and angular distribution in the altitude region between 120 and 1000 km are presented. The distribution is found to be isotropic at all altitudes below 250 km, while above this altitude anisotropies in both pitch angle and energy are found. The isotropy found in the angular distribution below 250 km implies that photoelectron transport below 250 km is insignificant, while the angular anisotropy found above this altitude implies a net photoelectron current in the upward direction. The energy anisotropy above 500 km arises from the selective backscattering of the low energy photoelectron population of the upward flux component by Coulomb collisions with the ambient ions. The total photoelectron flux attains its maximum value between about 40 and 70 km above the altitude at which the photoelectron production rate is maximum. The displacement of the maximum of the equilibrium flux is attributed to an increasing (with altitude) photoelectron lifetime. Photoelectrons at altitudes above that where the flux is maximum are on the average more energetic than those below that altitude. The flux of photoelectrons escaping to the protonosphere at dawn was found to be 2.6 × 10⁸ cm^?2 sec^?1, while the escaping flux at noon was found to be 1.5 × 10⁸ cm^?2 sec^?1. The corresponding escaping energy fluxes are: 4.4 × 10⁹ eV cm^?2 sec^?1 and 2.7 × 10⁹ eV cm^?2 sec^?1.     

Experimental study on the velocity of fragments in collisional breakup

The motion of fragments following a catastrophic destruction by either a normal or an oblique impact at 2.5–2.9 km sec^?1 into cubic and spherical basalt targets was studied with a high-speed framing camera. Velocities at the antipodes of the targets vary as (E/M)^0.75 (E = impact energy; M = target mass) and are lower than 200 m sec^?1 at E/M ? 10⁹ ergs g^?1. Excluding fine-grained particles from the impact site, 70 to 80% by mass fraction of the fragments have velocities lower than twice the antipodal velocity. Comminution and ejection energies wasted in this mass fraction were a few percent of the impact energy at E/M ? 5 × 10⁷ ergs g^?1. During a catastrophic impact into asteroids some of the fragmented bodies can be reconcentrated by mutual gravitation.     

Regolith history of lunar meteorites

Abstract— The regolith evolution of the lunar meteorites Dhofar (Dho) 081, Northwest Africa (NWA) 032, NWA 482, NWA 773, Sayh al Uhaymir (SaU) 169, and Yamato (Y‐) 981031 was investigated by measuring the light noble gases He, Ne, and Ar. The presence of trapped solar neon in Dho 081, NWA 773, and Y‐981031 indicates an exposure at the lunar surface. A neon three‐isotope diagram for lunar meteorites yields an average solar ²⁰Ne/²²Ne ratio of 12.48 ± 0.07 representing a mixture of solar energetic particles neon at a ratio of 11.2 and solar wind neon at a ratio of 13.8. Based on the production rate ratio of ²¹Ne and ³⁸Ar, the shielding depth in the lunar regolith of NWA 032, NWA 482, SaU 169, and Y‐981031 was obtained. The shielding depth of these samples was between 10.5 g/cm² and >500 g/cm². Based on spallogenic Kr and Xe, the shielding depth of Dho 081 was estimated to be most likely between 120 and 180 g/cm². Assuming a mean density of the lunar regolith of 1.8 g/cm³, 10.5 g/cm² corresponds to a depth of 5.8 cm and 500 g/cm² to 280 cm below the lunar surface. The range of regolith residence time observed in this study is 100 Ma up to 2070 Ma.