Study of Ribbon Separation and Magnetic Reconnection Rates in a Two-Ribbon Flare

In this paper, we present a study of the correlation between the speed of flare ribbon separation and the magnetic flux density during the 10 April 2001 solar flare. The study includes the section of the neutral line containing the flare core and its peripheral area. This event shows clear two-ribbon structure and inhomogeneous magnetic fields along the ribbons, so the spatial correlation and distribution of the flare and magnetic parameters can be studied. A weak negative correlation is found between the ribbon separation speed (V _r) and the longitudinal magnetic flux density (B _z ). This correlation is the weakest around the peak of the flare. Spatially, the correlation is also weakest at the positions of the hard X-ray (HXR) sources. In addition, we estimate the magnetic reconnection rate (electric field strength in the reconnection region E _rec) by combining the speed of flare ribbons and the longitudinal magnetic flux density. During flare evolution, the time profiles of the magnetic reconnection rate are similar to that of the ribbon separation speed, and the speeds of ribbon separation are relatively slow in the strong magnetic fields (i.e., V _r is negatively correlated with B _z ). However, along the flare ribbons, E _rec fluctuates in a small range except near the HXR source. A localized enhancement of the reconnection rate corresponds to the position of the HXR source.     

The Magnetic Structure and Generation of EUV Flare Ribbons

The `ribbons' of two-ribbon flares show complicated patterns reflecting the linkages of coronal magnetic field lines through the lower solar atmosphere. We describe the morphology of the EUV ribbons of the July 14, 2000 flare, as seen in SOHO, TRACE, and Yohkoh data, from this point of view. A successful co-alignment of the TRACE, SOHO/MDI and Yohkoh/HXT data has allowed us to locate the EUV ribbon positions on the underlying field to within ∼ 2′′, and thus to investigate the relationship between the ribbons and the field, and also the sites of electron precipitation. We have also made a determination of the longitudinal magnetic flux involved in the flare reconnection event, an important parameter in flare energetic considerations. There are several respects in which the observations differ from what would be expected in the commonly-adopted models for flares. Firstly, the flare ribbons differ in fine structure from the (line-of-sight) magnetic field patterns underlying them, apparently propagating through regions of very weak and probably mixed polarity. Secondly, the ribbons split or bifurcate. Thirdly, the amount of line-of-sight flux passed over by the ribbons in the negative and positive fields is not equal. Fourthly, the strongest hard X-ray sources are observed to originate in stronger field regions. Based on a comparison between HXT and EUV time-profiles we suggest that emission in the EUV ribbons is caused by electron bombardment of the lower atmosphere, supporting the hypothesis that flare ribbons map out the chromospheric footpoints of magnetic field lines newly linked by reconnection. We describe the interpretation of our observations within the standard model, and the implications for the distribution of magnetic fields in this active region.     

The formation of flare loops by magnetic reconnection and chromospheric ablation

Slow-mode shocks produced by reconnection in the corona can provide the thermal energy necessary to sustain flare loops for many hours. These slow shocks have a complex structure because strong thermal conduction along field lines dissociates the shocks into conduction fronts and isothermal subshocks. Heat conducted along field lines mapping from the subshocks to the chromosphere ablates chromospheric plasma and thereby creates the hot flare loops and associated flare ribbons. Here we combine a non-coplanar compressible reconnection theory with simple scaling arguments for ablation and radiative cooling, and predict average properties of hot and cool flare loops as a function of the coronal vector magnetic field. For a coronal field strength of 100 G the temperature of the hot flare loops decreases from 1.2 × 10⁷ K to 4.0 × 10⁶ K as the component of the coronal magnetic field perpendicular to the plane of the loops increases from 0% to 86% of the total field. When the perpendicular component exceeds 86% of the total field or when the altitude of the reconnection site exceeds 10⁶km, flare loops no longer occur. Shock enhanced radiative cooling triggers the formation of cool H flare loops with predicted densities of 10¹³ cm^–3, and a small gap of 10³ km is predicted to exist between the footpoints of the cool flare loops and the inner edges of the flare ribbons.     

Longitude variations of solar magnetic fields of different intensity in cycle 23 as inferred from the SOHO/MDI data

SOHO/MDI magnetograms have been used to analyze the longitude distribution of the squared solar magnetic field 〈B ²〉 in the activity cycle no. 23. The energy of the magnetic field (〈B ²〉) is shown to change with longitude. However, these variations hardly fit the concept of active longitudes. In the epochs of high solar activity, one can readily see a relationship between longitude variations of the medium-strong ((|B| > 50 G or |B| > 100 G) and relatively weak (|B| ≤ 50 G or |B| ≤ 100 G) fields at all latitudes. In other periods, this relationship is revealed mainly at the latitudes not higher than 30°. The background fields (|B| ≤ 25 G) also display longitude variations, which are, however, not related to those of the strong fields. This makes us think that the fields of solar activity are rather inclusions to the general field than the source of the latter.     

Average Thickness of Magnetosheath Upstream of Magnetic Clouds at 1 AU versus Solar Longitude of Source

Starting with a large number (N=100) of Wind magnetic clouds (MCs) and applying necessary restrictions, we find a proper set of N=29 to investigate the average ecliptic plane projection of the upstream magnetosheath thickness as a function of the longitude of the solar source of the MCs, for those cases of MCs having upstream shock waves. A few of the obvious restrictions on the full set of MCs are the need for there to exist a driven upstream shock wave, knowledge of the MC’s solar source, and restriction to only MCs of low axial latitudes. The analysis required splitting this set into two subsets according to average magnetosheath speed: slow/average (300 – 500 km s⁻¹) and fast (500 – 1100 km s⁻¹) speeds. Only the fast set gives plausible results, where the estimated magnetosheath thickness (ΔS) goes from 0.042 to 0.079 AU (at 1 AU) over the longitude sector of 0° (adjusted source-center longitude of the average magnetic cloud) to 40° off center (East or West), based on N=11 appropriate cases. These estimates are well determined with a sigma (σ) for the fit of 0.0055 AU, where σ is effectively the same as (chi-squared) for the appropriate quadratic fit. The associated linear correlation coefficient for ΔS versus |Longitude| was very good (c.c.=0.93) for the fast range, and ΔS at 60° longitude is extrapolated to be 2.7 times the value at 0°. For the slower speeds we obtain the surprising result that ΔS is typically more-or-less constant at 0.040±0.013 AU at all longitudes, indicating that the MC as a driver, when moving close to the normal solar wind speed, has little influence on magnetosheath thickness. In some cases, the correct choice between two candidate solar-source longitudes for a fast MC might be made by noting the value of the observed ΔS just upstream of the MC. Also, we point out that, for the 29 events, the average sheath speed was well correlated with the quantity ΔV[=(〈V _MC〉−〈V _UPSTREAM〉)], and also with both 〈V _MC〉 and 〈V _MC,T〉, where 〈V _MC〉 is the first one-hour average of the MC speed, 〈V _MC,T〉 is the average MC speed across the full MC, and 〈V _UPSTREAM〉 is a five-hour average of the solar wind speed just upstream of the shock.     

On a possible control mechanism for solar wind outflow velocity from coronal holes

In this paper we give an explanation for a control mechanism for velocityV of solar wind (SW) streams for coronal holes (CHs) based on the idea suggested by Rudenko and Fainshtein (1993). In accordance with that idea, the difference of values ofV in high-speed SW streams from different CHs is due to the spread in magnitude of magnetic fieldB _a in the region of acceleration of such streams near the Sun. In this case, with increasing magnitude ofB _a, there is an increase in velocity of the high-speed stream.Through calculations of the coronal magnetic field (potential-field approximation) it is shown that on the source surface the magnetic field B _s, averaged over the cross-section of the magnetic tube from a CH, can vary for different tubes over a wide range and correlates quite well with the area of this tube's base as well as depending on the radial component of the magnetic field at the base of the tube on the source surface B _or.It is found that the value of superradial divergence of the magnetic tube from a CH depends not only on the area of its base (as shown in prior work) but also on B _or. A positive correlation at the Earth's orbit between velocityV of the high-speed SW and the radial component of the magnetic field in the region of this stream is detected, which agrees indirectly with theV-control mechanism under discussion.     

A current ribbon model for energy storage and release with application to the flare of 7 January 1992

Observations of the flare on 7 January 1992 are interpreted using a topological model of the magnetic field. The model, developed here, applies a theory of three-dimensional reconnection to the inferred magnetic field configuration for 7 January. In the model field a new bipole ( 10²¹ Mx) emerges amidst pre-existing active region flux. This emergence gives rise to two current ribbons along the boundaries (separators) separating the distinct, new and old, flux systems. Sudden reconnection across these boundary curves transfers 3 ×10²⁰ Mx of flux from the bipole into the surrounding flux. The model also predicts the simultaneous (sympathetic) flaring of the two current ribbons. This explains the complex two-loop structure noted in previous observations of this flare. We subject the model predictions to comparisons with observations of the flare. The locations of current ribbons in the model correspond closely with those of observed soft X-ray loops. In addition the footpoints and apexes of the ribbons correspond with observed sources of microwave and hard X-ray emission. The magnitude of energy stored by the current ribbons compares favorably to the inferred energy content of accelerated electrons in the flare.     

Motions of Flare Ribbons and Loops in Various Magnetic Configurations

Kopp–Pneuman-type magnetic configurations, which include a vertical current sheet, with various background fields are investigated. Dissipation of the current sheet as a result of magnetic reconnection produces bright flare ribbons on the solar disk and a growing flare loop system in the corona. In principle, the growth of flare loop system is governed by a reconnection process only, and the behavior of flare ribbons is also controlled by the background field. The flare ribbons may appear either separate or attached to one another at the onset of the flare depending on the background field distribution on the boundary surface. We calculate the decrease in height that magnetic field lines undergo after they have reconnected to form closed loops. Following previous practice, we refer to this decrease as field line shrinkage. Unlike the motions of flare ribbons, the shrinkage of flare loops depends weakly on the background field. Individual loops always shrink fastest at the moment it is produced by reconnection and just starts to leave the current sheet. The earlier the loop forms, the more and faster it shrinks. The relevant observations are explained on the basis of our calculations, and the aspects of the explanation that need improvement are also discussed.     

Conservation of both current and helicity in a quadrupolar model for solar flares

A model for a solar flare, involving magnetic reconnection transferring flux and current between current-carrying magnetic loops connecting two pairs of footpoints, is generalized to include conservation of magnetic helicity during reconnection, as well as conservation of current at all four footpoints. For a set of force-free loops, with the ith loop having flux F _i and current I _i, the self and mutual helicities are proportional to the self and mutual inductances with the constant of proportionality determined by α_i=F _i/μ₀ I _i. In a constant-α model, the change in magnetic energy is proportional to the change in helicity, and conservation of helicity implies conservation of magnetic energy, so that a flare cannot occur. In a quadrupolar model, with α₁>α₂ initially, α₁ increases and α₂ decreases when flux and current are transferred from loops 1 and 2 to loops 3 and 4. A model that conserves both current and helicity is constructed; it depends on the initial αs, and otherwise is somewhat simpler than when helicity is neglected.     

Evidence for Tether-Cutting Reconnection in a Quadrupole Magnetic Configuration in the April 9, 2001, M7.9 Flare

We studied the M7.9 flare on April 9, 2001 that occurred within a δ-sunspot of active region NOAA 9415. We used a multi-wavelength data set, which includes Yohkoh, TRACE, SOHO, and ACE spacecraft observations, Potsdam and Ondřejov radio data and Big Bear Solar Observatory (BBSO) images in order to study the large-scale structure of this two-ribbon flare that was accompanied by a very fast coronal mass ejection (CME). We analyzed light curves of the flare emission as well as the structure of the radio emission and report the following: the timing of the event, i.e., the fact that the initial brightenings, associated with the core magnetic field, occurred earlier than the remote brightening (RB), argue against the break-out model in the early phase of this event. We thus conclude that the M7.9 flare and the CME were triggered by a tether-cutting reconnection deep in the core field connecting the δ-spot and this reconnection formed an unstable flux rope. Further evolution of the erupted flux rope could be described either by the “standard“ flare model or a break-out type of the reconnection. The complex structure of flare emission in visible, X-ray, and radio spectral ranges point toward a scenario which involves multiple reconnection processes between extended closed magnetic structures.     

Homologous Flares and Magnetic Field Topology in Active Region NOAA 10501 on 20 November 2003

We present and interpret observations of two morphologically homologous flares that occurred in active region (AR) NOAA 10501 on 20 November 2003. Both flares displayed four homologous Hα ribbons and were both accompanied by coronal mass ejections (CMEs). The central flare ribbons were located at the site of an emerging bipole in the centre of the active region. The negative polarity of this bipole fragmented in two main pieces, one rotating around the positive polarity by ≈ 110° within 32 hours. We model the coronal magnetic field and compute its topology, using as boundary condition the magnetogram closest in time to each flare. In particular, we calculate the location of quasi-separatrix layers (QSLs) in order to understand the connectivity between the flare ribbons. Though several polarities were present in AR 10501, the global magnetic field topology corresponds to a quadrupolar magnetic field distribution without magnetic null points. For both flares, the photospheric traces of QSLs are similar and match well the locations of the four Hα ribbons. This globally unchanged topology and the continuous shearing by the rotating bipole are two key factors responsible for the flare homology. However, our analyses also indicate that different magnetic connectivity domains of the quadrupolar configuration become unstable during each flare, so that magnetic reconnection proceeds differently in both events.     

Multiwavelength Observations of a Failed Flux Rope in the Eruption and Associated M-Class Flare from NOAA AR 11045

We present the multiwavelength observations of a flux rope that was trying to erupt from NOAA AR 11045 and the associated M-class solar flare on 12 February 2010 using space-based and ground-based observations from TRACE, STEREO, SOHO/MDI, Hinode/XRT, and BBSO. While the flux rope was rising from the active region, an M1.1/2F class flare was triggered near one of its footpoints. We suggest that the flare triggering was due to the reconnection of a rising flux rope with the surrounding low-lying magnetic loops. The flux rope reached a projected height of ≈0.15R _⊙ with a speed of ≈90 km s⁻¹ while the soft X-ray flux enhanced gradually during its rise. The flux rope was suppressed by an overlying field, and the filled plasma moved towards the negative polarity field to the west of its activation site. We found the first observational evidence of the initial suppression of a flux rope due to a remnant filament visible both at chromospheric and coronal temperatures that evolved a couple of days earlier at the same location in the active region. SOHO/MDI magnetograms show the emergence of a bipole ≈12 h prior to the flare initiation. The emerged negative polarity moved towards the flux rope activation site, and flare triggering near the photospheric polarity inversion line (PIL) took place. The motion of the negative polarity region towards the PIL helped in the build-up of magnetic energy at the flare and flux rope activation site. This study provides unique observational evidence of a rising flux rope that failed to erupt due to a remnant filament and overlying magnetic field, as well as associated triggering of an M-class flare.     

Ion-Sound Model of Microwave Spikes with Fast Shocks in the Reconnection Region

A new model for solar spike bursts is considered based on the interaction of Langmuir waves with ion-sound waves: l+st. Such a mechanism can operate in shock fronts, propagating from a magnetic reconnection region. New observations of microwave millisecond spikes are discussed. They have been observed in two events: 4 November 1997 between 05:52–06:10 UT and 28 November 1997 between 05:00–05:10 UT using the multichannel spectrograph in the range 2.6–3.8 GHz of Beijing AO. Yohkoh/SXT images in the AR and SOHO EIT images testify to a reconstruction of bright loops after the escape of a CME. A fast shock front might be manifested as a very bright line in T _e SXT maps (up to 20 MK) above dense structures in emission measure (EM) maps. Moreover one can see at the moment of spike emission (for the 28 November 1997 event) an additional maximum at the loop top on the HXR map in the AR as principal evidence of fast shock propagation. The model gives the ordinary mode of spike emission. Sometimes we observed a different polarization of microwave spikes that might be connected with the depolarization of the emission in the transverse magnetic field and rather in the vanishing magnetic field in the middle of the QT region. Duration and frequency band of isolated spikes are connected with parameters of fast particle beams and shock front. Millisecond microwave spikes are probably a unique manifestation of flare fast shocks in the radio emission.     

The Evaporation Regime in a Confined Flare

We studied the evolution of a small eruptive flare (GOES class C1) from its onset phase using multi-wavelength observations that sample the flare atmosphere from the chromosphere to the corona. The main instruments involved were the Coronal Diagnostic Spectrometer (CDS) aboard SOHO and facilities at the Dunn Solar Tower of the National Solar Observatory/Sacramento Peak. Transition Region and Coronal Explorer (TRACE) together with Ramaty High-Energy Spectroscopic Imager (RHESSI) also provided images and spectra for this flare. Hα and TRACE images display two loop systems that outline the pre-reconnection and post-reconnection magnetic field lines and their topological changes revealing that we are dealing with an eruptive confined flare. RHESSI data do not record any detectable emission at energies ≥25 keV, and the observed count spectrum can be well fitted with a thermal plus a non-thermal model of the photon spectrum. A non-thermal electron flux F ≈ 5 × 10¹⁰ erg cm⁻² s⁻¹ is determined. The reconstructed images show a very compact source whose peak emission moves along the photospheric magnetic inversion line during the flare. This is probably related to the motion of the reconnection site, hinting at an arcade of small loops that brightens successively. The analysis of the chromospheric spectra (Ca II K, He I D₃ and Hγ, acquired with a four-second temporal cadence) shows the presence of a downward velocity (between 10 and 20 km s⁻¹) in a small region intersected by the spectrograph slit. The region is included in an area that, at the time of the maximum X-ray emission, shows upward motions at transition region (TR) and coronal levels. For the He I 58.4 and O v 62.97 lines, we determine a velocity of ≈−40 km s⁻¹ while for the Fe XIX 59.22 line a velocity of ≈−80 km s⁻¹ is determined with a two-component fitting. The observations are discussed in the framework of available hydrodynamic simulations and they are consistent with the scenario outlined by Fisher (1989). No explosive evaporation is expected for a non-thermal electron beam of the observed characteristics, and no gentle evaporation is allowed without upward chromospheric motion. It is suggested that the energy of non-thermal electrons can be dissipated to heat the high-density plasma, where possibly the reconnection occurs. The consequent conductive flux drives the evaporation process in a regime that we can call sub-explosive.     

Two Regularities in the Coronal Green-Line Brightness – Magnetic Field Coupling and the Heating of the Corona

To study the quantitative relationship between the brightness of the coronal green line 530.5 nm Fe xiv and the strength of the magnetic field in the corona, we have calculated the cross-correlation of the corresponding synoptic maps for the period 1977 – 2001. The maps of distribution of the green-line brightness I were plotted using every-day monitoring data. The maps of the magnetic field strength B and the tangential B _t and radial B _r field components at the distance 1.1 R _⊙ were calculated under potential approximation from the Wilcox Solar Observatory (WSO) photospheric data. It is shown that the correlation I with the field and its components calculated separately for the sunspot formation zone ±30° and the zone 40 – 70° has a cyclic character, the corresponding correlation coefficients in these zones changing in anti-phase. In the sunspot formation zone, all three coefficients are positive and have the greatest values near the cycle minimum decreasing significantly by the maximum. Above 40°, the coefficients are alternating in sign and reach the greatest positive values at the maximum and the greatest negative values, at the minimum of the cycle. It is inferred that the green-line emission in the zone ±30° is mainly controlled by B _t, probably due to the existence of low arch systems. In the high-latitude zone, particularly at the minimum of the cycle, an essential influence is exerted by B _r, which may be a manifestation of the dominant role of large-scale magnetic fields. Near the activity minimum, when the magnetic field organization is relatively simple, the relation between I and B for the two latitudinal zones under consideration can be represented as a power-law function of the type I ∝ B ^q. In the sunspot formation zone, the power index q is positive and varies from 0.75 to 1.00. In the zone 40 – 70°, it is negative and varies from −0.6 to −0.8. It is found that there is a short time interval approximately at the middle of the ascending branch of the cycle, when the relationship between I and B vanishes. The results obtained are considered in relation to various mechanisms of the corona heating.     

Predictions of Energy and Helicity in Four Major Eruptive Solar Flares

In order to better understand the solar genesis of interplanetary magnetic clouds (MCs), we model the magnetic and topological properties of four large eruptive solar flares and relate them to observations. We use the three-dimensional Minimum Current Corona model (Longcope, 1996, Solar Phys. 169, 91) and observations of pre-flare photospheric magnetic field and flare ribbons to derive values of reconnected magnetic flux, flare energy, flux rope helicity, and orientation of the flux-rope poloidal field. We compare model predictions of those quantities to flare and MC observations, and within the estimated uncertainties of the methods used find the following: The predicted model reconnection fluxes are equal to or lower than the reconnection fluxes inferred from the observed ribbon motions. Both observed and model reconnection fluxes match the MC poloidal fluxes. The predicted flux-rope helicities match the MC helicities. The predicted free energies lie between the observed energies and the estimated total flare luminosities. The direction of the leading edge of the MC’s poloidal field is aligned with the poloidal field of the flux rope in the AR rather than the global dipole field. These findings compel us to believe that magnetic clouds associated with these four solar flares are formed by low-corona magnetic reconnection during the eruption, rather than eruption of pre-existing structures in the corona or formation in the upper corona with participation of the global magnetic field. We also note that since all four flares occurred in active regions without significant pre-flare flux emergence and cancelation, the energy and helicity that we find are stored by shearing and rotating motions, which are sufficient to account for the observed radiative flare energy and MC helicity.     

Flare Plasma Cooling from 30 MK down to 1 MK modeled from Yohkoh,GOES, and TRACE observations during the Bastille Day Event (14 July 2000)

We present an analysis of the evolution of the thermal flare plasma during the 14 July 2000, 10 UT, Bastille Day flare event, using spacecraft data from Yohkoh/HXT, Yohkoh/SXT, GOES, and TRACE. The spatial structure of this double-ribbon flare consists of a curved arcade with some 100 post-flare loops which brighten up in a sequential manner from highly-sheared low-lying to less-sheared higher-lying bipolar loops. We reconstruct an instrument-combined, average differential emission measure distribution dEM(T)/dT that ranges from T=1 MK to 40 MK and peaks at T ₀=10.9 MK. We find that the time profiles of the different instrument fluxes peak sequentially over 7 minutes with decreasing temperatures from T≈30 MK to 1 MK, indicating the systematic cooling of the flare plasma. From these temperature-dependent relative peak times t _peak(T) we reconstruct the average plasma cooling function T(t) for loops observed near the flare peak time, and find that their temperature decrease is initially controlled by conductive cooling during the first 188 s, T(t)∼[1+(t/τ_cond)]^−2/7, and then by radiative cooling during the next 592 s, T(t)∼[1−(t/τ_rad)]^3/5. From the radiative cooling phase we infer an average electron density of n _e=4.2×10¹¹ cm⁻³, which implies a filling factor near 100% for the brightest observed 23 loops with diameters of ∼1.8 Mm that appear simultaneously over the flare peak time and are fully resolved with TRACE. We reproduce the time delays and fluxes of the observed time profiles near the flare peak self-consistently with a forward-fitting method of a fully analytical model. The total integrated thermal energy of this flare amounts to E _thermal=2.6×10³¹ erg. Supplementary material to this paper is available in electronic form at http://dx.doi.org/10.1023/A:1014257826116     

Some results of the cooperative photometric observations of the UV Cet-type flare stars in the years 1967–71

The list of the cooperative photometric observations of the UV Cet-type flare stars that have been organized during the years 1967 to 1971 by the Working Group on Flare Stars of the IAU Commission 27 is given. The completeness and reliability of the data obtained are evaluated by comparing simultaneous observations at different observatories. the statistical analysis of the UV Cet, YZ CMi, EV Lac and AD Leo flares observed in the B-band is carried out. The flare energy spectrum in the energy range where observational selection effects are small is found to be d lnv/d lnE _B=–1.4 to –1.9,v is an occurrence of flares with radiation energy ofE _B. The total flares' radiation is equal to 1.7%, 1.2%, 0.3% and 0.4% of the quiet radiation in the B-band of the stars listed, and the main part of this total radiation is due to the strongest flares. Distributions of flare rise times (t _r) and of rates of flare absolute luminosity increase (d² E _B/dt _r ²Ë _r) are considered; these parameters of flare are independent statistically for all stars studied. Correlation coefficientsr (E _B,t _r) andr (E _B,r(E _B,Ë _r)) are rather small except r (E _B,t _r)=0.86 for the AD Leo flares. Contradictory conclusions on temporal distribution of flares infered by different investigators are noted.     

Analysis of Magnetic Shear in An X17 Solar Flare on October 28, 2003

An X17 class (GOES soft X-ray) two-ribbon solar flare on October 28, 2003 is analyzed in order to determine the relationship between the timing of the impulsive phase of the flare and the magnetic shear change in the flaring region. EUV observations made by the Transition Region and Coronal Explorer (TRACE) show a clear decrease in the shear of the flare footpoints during the flare. The shear change stopped in the middle of the impulsive phase. The observations are interpreted in terms of the splitting of the sheared envelope field of the greatly sheared core rope during the early phase of the flare. We have also investigated the temporal correlation between the EUV emission from the brightenings observed by TRACE and the hard X-ray (HXR) emission (E > 150 keV) observed by the anticoincidence system (ACS) of the spectrometer SPI on board the ESA INTEGRAL satellite. The correlation between these two emissions is very good, and the HXR sources (RHESSI) late in the flare are located within the two EUV ribbons. These observations are favorable to the explanation that the EUV brightenings mainly result from direct bombardment of the atmosphere by the energetic particles accelerated at the reconnection site, as does the HXR emission. However, if there is a high temperature (T > 20 MK) HXR source close to the loop top, a contribution of thermal conduction to the EUV brightenings cannot be ruled out.     

The Gradual Phase of the X17 Flare on October 28, 2003

We present SOHO/CDS observations taken during the gradual phase of the X17 flare that occurred on October 28, 2003. The CDS data are supplemented with TRACE and ground-based observations. The spectral observations allow us to determine velocities from the Doppler shifts measured in the flare loops and in the two ribbon kernels, one hour and a half after the flare peak. Strong downflows (>70 km s⁻¹) are observed along the loop legs at transition-region temperatures. The velocities are close to those expected for free fall. Observations and results from a hydrodynamic simulation are consistent with the heating taking place for a short time near the top of the arcade. Slight upflows are observed in the outer edges of the ribbons (<60 km s⁻¹) in the EUV lines formed at log T < 6.3. These flows could correspond to the so-called “gentle evaporation.” At “flare” temperatures (Fe xix, log T = 6.9), no appreciable flows are observed. The observations are consistent with the general standard reconnection models for two-ribbons flares.