MULTIWAVELENGTH ANALYSIS OF THE IMPACT POLARIZATION OF 2001 JUNE 15 SOLAR FLARE

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1 The Astrophysical Journal, 631: , 2005 September 20 # The American Astronomical Society. All rights reserved. Printed in U.S.A. MULTIWAVELENGTH ANALYSIS OF THE IMPACT POLARIZATION OF 2001 JUNE 15 SOLAR FLARE Z. Xu 1 Laboratoire d Etudes Spatiales et d Instrumentation en Astrophysique, Observatoire de Paris, 5 Place Jules Janssen, F Meudon Cedex, France; and Astronomy Department, University of Nanjing, Nanjing, China J.-C. Hénoux and G. Chambe Laboratoire d Etudes Spatiales et d Instrumentation en Astrophysique, Observatoire de Paris, 5 Place Jules Janssen, F Meudon Cedex, France M. Karlický Astronomical Institute, Academy of Sciences of the Czech Republic, CZ Ond rejov, Czech Republic and C. Fang Astronomy Department, University of Nanjing, Nanjing, China Received 2005 March 20; accepted 2005 May 28 ABSTRACT We report here on the temporal and spatial evolution of the impact polarization of the H and H lines during an M6.3 solar flare observed on 2001 June 15 with the THEMIS telescope in the multiwavelength spectropolarimetric mode. Typical spectral intensity and polarization profiles are presented. Both lines are linearly polarized. The H line degree of polarization exceeds 4% at line center and in the near line wings. The H line is also linearly polarized, with a degree of polarization reaching 6%. The directions of polarization are either parallel or perpendicular to the local transverse magnetic field (i.e., either radial or tangential because the transverse magnetic field is directed almost in the flare-to-disk center direction). However, contrary to H,theH polarization direction is radial only. The H and H polarization islands are located at the edges of flare kernels. Only for radial polarization are these islands cospatial. No H polarization is found at the places where tangential H polarization is present. The origin of the observed polarization is discussed. Bombardment by low-energy protons or high-energy electrons associated with return currents can explain the radial polarization observed in the lowest flare kernel. The tangential H polarization observed in the surge near the upper flare location is interpreted as due to the electric current at the origin of the electromagnetic force that lifts the surge. Subject headings: polarization Sun: flares 1. INTRODUCTION During the rise or maximum phase of soft X-ray emission, the chromospheric H and H lines have been found to be linearly polarized in spectroscopic observations and filtergrams made during a few solar flares. H polarization was reported by Hénoux & Semel (1981), Vogt & Hénoux (1996, 1999), Vogt et al. (2001), and Hanaoka (2003) in two-dimensional filtergram observations covering flaring active regions. The linear polarization was directed along the flare toward the disk center with a degree of polarization 1% 5%. Independently, spectroscopic observations made with THEMIS (French-Italian solar telescope) and the Large Solar Vacuum Telescope (Baikal Astrophysical Observatory of Russia) confirmed the presence of linear polarization during flares (Firstova & Kashapova 2002). In particular, multiwavelength observations made in the H and H lines, which provide information on the variation of the polarization with depth ( Firstova et al. 1997; Hénoux & Karlický 2003), showed a 90 change in the linear polarization direction from tangential in H to radial in H (Hénoux & Karlický 2003). It has been recognized in atomic physics that an anisotropic velocity distribution of fast charged particles gives rise, via impact excitation, to an alignment called self-alignment 1 Send offprint requests to Z. Xu: xuzhi@nju.edu.cn. 618 (Kazantsev & Hénoux 1995) in an ensemble of atoms, ions, and molecules. Consequently, the radiation emitted could be polarized and could have an anisotropic intensity angular distribution. The highest degree of polarization is observed at 90 from the beam propagation direction. The direction of linear polarization and the degree of polarization depend on the energy of the incoming particle (Kleinpoppen 1969). In brief, the direction of polarization rotates by 90 at a turnover energy E tov,where the degree of polarization becomes null. In the case of the H line, for excitation by a monoenergetic, monodirectional electron or proton beam and for observations made at 90 of this beam, the degree of polarization reaches 30% at the 12 ev excitation threshold E thr with a polarization direction parallel to the particle beam velocity direction. The turnover energies E tov are 200 ev for electrons and 400 kev for protons (Werner & Schartner 1996). In H, the degree of polarization reaches 40% at threshold with a polarization direction also parallel to the particle beam velocity direction (A. Petreshen 2004, private communication). In the solar atmosphere, the particle velocity distribution is expected to be symmetrical around the magnetic field. Protons and electrons of energy below E tov, with a null pitch angle will produce a linear polarization directed along the magnetic field projection on the solar disk, i.e., radial for a vertical magnetic field. High-velocity electrons or protons of energy above E tov will lead to a direction of polarization perpendicular to this

2 IMPACT POLARIZATION OF FLARE 619 projection. On the other hand, for electrons or protons with a 90 pitch angle circling around the magnetic field, these conclusions must be reversed. The linear polarization of the H and H lines observed in solar flares has been interpreted as impact polarization due to collisional excitation of hydrogen by particles with an anisotropic velocity distribution. However, the nature and energy of these particles and the origin of their anisotropy are still debated. Low-velocity protons have been considered to be good candidates (Hénoux et al. 1990; Fletcher & Brown 1998), especially when there is no simultaneous hard X-ray or microwave impulsive emission. Another argument in favor of protons being at the origin of the observed polarization comes from the excessive energy requirement resulting from the long duration of impact polarization (Fletcher & Brown 1995) in case it were due to high-energy electrons. Low-energy electrons are easily scattered and their velocity distribution made isotropic in collisions. Therefore, impact polarization by low-energy electrons requires the anisotropy of their velocity distribution function to be produced or driven locally. A strong temperature gradient between the corona and chromosphere was proposed by Aboudarham et al. (1992). More recently, collisional excitation by the electrons of the return current, which neutralizes the electron beam current during electron bombardment, was proposed ( Karlický & Hénoux 2002; Hénoux & Karlický 2003). Impact line-linear polarization is very sensitive to particle propagation conditions such as local density, coronal mass, and particle nature and energy. Therefore, impact polarization is expected not to be present all over the flaring region but rather at locations where kinetic particles still have anisotropic velocity distributions after crossing the corona and where the local Fig. 1. Orientation reference axis, defining the Q and U Stokes parameters given by THEMIS. There exists a 45 angle between the celestial north-south direction OS and the OX axis that defines with OY the Stokes parameter Q (Q ¼ I X I Y,whereI X and I Y are, respectively, the line intensities along OX and OY). The cross locates the flare of 2001 June 15, and the dashed line points out the disk center-to-flare direction, which forms with the reference axis OS an angle close to 65. Fig. 2. Time evolution of soft X-ray emission (1 8 and ) for the 2001 June 15 flare. THEMIS observations began at 10:07 UT, as indicated by the vertical line. hydrogen density is not too high in order to avoid strong depolarization effects (Vogt et al. 2001). In this paper, the time and space evolution of the H and H line impact polarization during the 2001 June 15 event, reported in Hénoux & Karlický (2003), is studied in more detail. The main features of the THEMIS observations are introduced in x 2. In x 3 we discuss the properties of the two-dimensional spatial distribution, its relationship with the magnetic field and velocity field, and the linear polarization wavelength dependence. A general discussion and a conclusion are given in x THEMIS FLARE SPECTROPOLARIMETRIC OBSERVATIONS On 2001 June 15, an M6.3 flare (Geostationary Operational Environmental Satellite [GOES] 8 soft X-ray class) beginning at 10:01 UT in the NOAA Active Region 9502, and accompanied by a powerful mass ejection, was observed from10:07:20 to 10:30:56 UT by THEMIS in the multiraise (MTR) multiline spectropolarimetric mode. 2 In this mode the Stokes parameters of a set of lines are measured simultaneously. Figure 1 locates the flare (S26 E41 ) and shows the orientations of the reference axis used to define the Q and U Stokes parameters. The celestial north-south direction is the reference axis for the orientation of the linear polarization. The flare to disk center direction forms an angle close to 65 with this axis. In the hypothesis that the local magnetic field is close to vertical, the solar flare was located far enough from the solar disk center to give observation conditions favorable to impact polarization measurements. The flare region was scanned repeatedly 16 times with a 1 00 width spectrograph entrance slit. Each scan was covered in 20 steps separated by 2B5 in space and 4.5 s in time. Along the slit the spatial separation of two consecutive pixels corresponds to 0B42 on the solar image. For each given position (step) of the slit on the solar image, two spectra formed by the extraordinary and ordinary beams carrying the I þ S and I S signals, respectively, where S is either Q or U, were simultaneously recorded for a set of lines including H,H,Fei ( ), Na i D1 and D2; Q and U are measured alternatively with a 2.3 s time interval. The flare was associated with soft X-ray emission as measured by GOES. The temporal behavior of this soft X-ray emissionisshowninfigure2.hardx-rayandradioflux impulsive emissions were also present (Figs. 3 and 4 of Hénoux & Karlický 2003), which usually indicates that electron beams are bombarding the chromosphere. We focus here on the observations made from 10:07 to 10:15 UT, when strong polarization signals were present. 2 See

3 620 XU ET AL. Vol LINEAR POLARIZATION ANALYSIS Recent impact polarization observations have led to the conclusion that (1) H linear polarization has been found at the edge of a major flare by using a fixed slit (Firstova & Xu 2003), (2) the polarization direction changes with position in the flaring region (Xu et al. 2003), (3) the decrease of the degree of polarization is associated with an increase of the line intensity and width (Hénoux & Karlický 2003) in flares, and (4) the polarization in flares and in surges has perpendicular directions (Firstova & Kashapova 2002) Deriving the State of Polarization of the Lines Observed All the raw data were normalized using flat-field and darkcurrent measurements. The dark-current time variation was taken into account as explained in Bommier & Rayrole (2002). First, the line-relative Stokes parameter S/I is obtained by the combinations of the I þ S and I S signals as S=I ¼ (I þ S ) (I S ) (I þ S ) þ (I S ) : ð1þ Therefore, in order to be sure to compare signals coming from the same location on the Sun, the relative positioning of the I þ S and I S spectra must be done with the highest possible accuracy. This positioning requires the correction of the data from line curvature and from a slight relative rotation and difference in size along the slit direction of the two spectra. The method used to obtain an accurate relative positioning is discussed in more detail in the Appendix. Then we move from the Stokes representation of the light polarization to the representation based on the degree of polarization P: p P ¼ ffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi U 2 þ Q 2 =I and the angle that defines the direction of polarization, i.e., the orientation of the electric induction vector: or ¼ 0:5arctan(U=Q) þ 45 for Q > 0; ¼ 0:5arctan(U=Q) þ for Q < 0: This representation is used in the forthcoming sections Spatial Distribution of the Degree and Orientation of the Polarization As seen from the Earth for charged particles moving along the magnetic field, the direction of impact polarization is either along the direction of the transverse field or perpendicular to it. Therefore, the impact polarization direction must be compared to the direction of the local transverse magnetic field. The twodimensional distribution of the orientation of the transverse field was deduced from the Q and U profiles of the Fe i ( )line. As shown in Figure 3, the transverse field does not deviate significantly from the flare to disk center direction (thick arrows). The time dependence of the H and H line center (0.02 nm) intensity I, polarization orientation, and degree of polarization P are shown in Figure 4 in two-dimensional maps of the flaring region. The beginning scan times for each twodimensional image are 10:07:20, 10:08:54, 10:10:29, 10:13:38, Fig. 3. Two-dimensional representation of the transverse magnetic field (dashes). The transverse field direction stays oriented close to the flare to disk center direction (arrow). H line center intensity contours at 10:07 UTare also shown. and 10:15:13 UT. Line-linear polarization is present in a few polarization islands, generally at the edges of the flare rather than in the flare kernels, with a degree of polarization much higher than the noise level (0.7%). The degree of polarization reaches 4% in H and 6% in H line centers in the early phase and then fades away with time. The positions along the scan direction of the linear polarization islands are indicated by arrows in Figure 4 and their ordinates x and y given in Table 1, where all the polarization characteristics of these islands and the orientations of the transverse magnetic field are listed. From Table 1 one first concludes that, as expected for a nearly vertical magnetic field, the transverse magnetic field is close to the flare to disk center direction with a deviation varying from 22 to 10. Second, there is a significant tendency for the H and H polarization directions to be parallel (indicated by k)or perpendicular (indicated by?) to the transverse magnetic field within 20 as reflected by the way the islands are denominated. Their names are made of three characters XYZ, where X and Y describe the states of H and H linear polarization (R for radial, T for tangential, N for no polarization, and? for random polarization directions), respectively, and Z differentiates the various polarization islands. The existence of different islands denominated by different values of Z may just be due to the time evolution of a single island. Discontinuities in the flare temporal coverage make it difficult to closely follow the time evolution of the polarization islands. However, the spatial evolution of the TNZ islands in relation to the spatial and time evolution of the velocity field suggests that TN1, TN2, and TN3 islands correspond to the same moving structure as discussed in the next subsections and visualized in Figure 8. During the whole flare observing time, both radial and tangential polarization directions are present in H.Contrary to H, the H line is polarized only radially. The majority of locations with radial polarization in H are cospatial with islands of H radial polarization (RR1 RR4). This is not the case for H tangential polarization islands ( TN1 TN3), where no significant H is present.

4 No. 1, 2005 IMPACT POLARIZATION OF FLARE 621 Fig. 4. Two-dimensional representation of H (top row) and H (bottom row) polarization degree and orientation superposed on their intensity maps (0B42 pixel 1 ), as a function of time (increasing from right to left). The beginning scan times are 10:07:20, 10:08:54, 10:10:29, 10:13:38, and 10:15:13 UT. Each scan is completed in about 90 s. The isocontours of polarization degree correspond, respectively, to P ¼ 2:25%, 3.3%, and 4.4% in H and to P ¼ 3:5%, 4.4%, and 5.0% in H line centers (0.02 nm). The polarization directions in the islands delimited by these contours are represented as white (gray) pixels for the radial (tangential) directions (within 15 ). Arrows point out the slit position in a scan crossing a polarization island usually at polarization peaks. The names of these parts of islands are made of three characters XYZ, where X and Y describe the states of H and H linear polarization (R for radial, T for tangential, N for no polarization, and? for random polarization directions), respectively, and Z differentiates the various cuts. Associated intensity and polarization profiles are drawn in Figs. 6 and 7. There is not a one-to-one relation between the polarization orientation in a given line and the polarity of the longitudinal field. Radial and tangential H polarization islands are not separated by the longitudinal magnetic neutral line (see Fig. 5). They are both observed in a single-polarity negative region of the magnetic field Polarization Wavelength Dependence The wavelength dependence of the line intensity, degree of polarization, and polarization direction of eight of the polarization islands present in Figure 4 is investigated in this section and illustrated in Figures 6 and 7. In these two figures, on the left of each set of profiles corresponding to a given polarization island the variations along the slit direction of both the degree of polarization and the intensity gradient amplitude are compared. This is done in order to verify that the relative positioning of the two I þ S and I S spectra was done properly (see the Appendix). The examples of the radial polarization illustrated in Figure 6 correspond to the four locations RR1 RR4 observed in both H and H lines. In these examples there is no evident correlation between the degree of polarization and the intensity gradient along the slit. Especially in examples 2, 3, 4, 7, and 8, the intensity gradient TABLE 1 Directions of the Impact Polarization and of the Transverse Magnetic Field Time (UT) Position ( pixels) Arrow Polarization Direction Angle of H (deg) Polarization Direction Angle of H (deg) B? (deg) 10: < x < 35, 55 < y < 65 RR (k) 82.5 (k) < x < 50, 65 < y < 80 RR (k) 70.9 (k) < x < 55, 20 < y < 30 a 96.1 No < x < 60, 70 < y < 80 TN (?) No : < x < 45, 60 < y < 70 TN (?) No : < x < 30, 60 < y < 70 NR1 No 69.1 (k) < x < 40, 55 < y < 75 TN (?) No : < x < 25, 65 < y < 75 RR (k) 63.3 (k) < x < 60, 80 < y < 85 a?r1? 64.7 (k) : < x < 10, 65 < y < 80 RR (k) 65.5 (k) < x < 40, 75 < y < 85 NR2 No 65.7 (k) 49.2 Notes. The mean direction angles of the polarization at the center of the H and H lines and of the mean transverse magnetic field direction are obtained by integrating over the areas delimited by the contour lines of the polarization degree represented in Fig. 4. The OS reference axis is represented in Fig. 1. The flare-to-disk direction makes a angle with OS. The pixel width corresponds to 0B42. a The area is close to a sunspot.

5 622 XU ET AL. Fig. 5. Comparison of the two-dimensional distributions of H (left) and H (right) polarization observed at 10:07 UT with a two-dimensional map longitudinal magnetic field. The field was derived from the Stokes V profiles of the Fe i ( ) line observed at 09:55 UT by using the center-of-gravity method. Solid line (dotted line) represents the positive (negative) magnetic field. The impact polarization characteristics are represented as in Fig. 4. is almost 0 where the polarization is significant. The majority of the H intensity I profiles show typical center-reversal emission, which is usually regarded as evidence of nonthermal particle bombardment on the chromosphere. The linear polarization is concentrated at line center with a remarkably constant direction. We now turn to the case of the tangential polarization observed in H in the TN1 to TN3 islands. The H intensity profiles shown in Figure 7 have a much deeper and broader absorption core than the ones plotted in Figure 6, and an obvious blueshift is present in both intensity and polarization degree profiles. This blueshift corresponds to an upward velocity close to 50 km s 1. The profiles show large macro- and microvelocities. Although a weak tangential polarization signal, at a 1 3 level (0.6% 2.4%), is observed in the red part of the H line profile in the TN1,TN2, and TN3 polarization islands, most of the tangential polarization signal is concentrated in a strong polarization peak in the blue part of the line. The blueshifted polarization reaches 4.3% 6.5%, and the larger the velocity disturbance, the stronger the polarization is. As shown in Figure 8 the location of the upflows detected in H moves north-south with a velocity of 70 km 1. The associated shift in position of the TN islands suggests that they correspond to a single polarization island that is moving with time with a surge. 4. DISCUSSION AND CONCLUSION A multiwavelength observation of a single flare in H and in hard X-ray (RHESSI) has shown that hard X-ray emission (due to electron bombardment) was present not on the brightest part of the H flare but rather at the periphery of the H kernels (Kašparová 2004). Accordingly, in the 2001 June flare the linear polarization of H and H lines, free from the effect of intensity gradient (see the Appendix), is present only at the edges of the flare kernels where it covers areas of a few times cm 2. For both H and H lines, the impact polarization observed perpendicularly to a monoenergetic beam has the same orientation near excitation threshold and rotates by 90 at high energies. Therefore, it is worth using these two lines to follow the beam penetration into the solar atmosphere. As expected, the observations show a tendency for the H and H polarization directions to be parallel or perpendicular, within 20,tothe transverse magnetic field direction deduced from the linear polarization of the Fe i ( ) line. Three main associations are present, as summarized in Table 1 and Figure 4: (1) radial direction of polarization in both H and H lines (RR1 to RR4), (2) tangential polarization in H without significant polarization in H (TN1 to TN3), and (3) radial polarization in H without significant polarization in H (NR1 to NR2). Assuming that the impacting particle velocity distribution functions are symmetrical along a nearly vertical magnetic field, we discuss below the possible origins of these associations. Radial polarization in two lines (RR). H is formed lower in the atmosphere than H, with a similar turnover energy E tov of about 200 ev (400 kev) for incident electron (proton) beams ( N. Feautrier 2002, private communication). The radial direction of polarization observed in both H and H lines with a good spatial correspondence requires the particles to reach the layers where the H and H lines are formed with an energy lower than E tov. The deceleration distance for 400 kev protons in a flaring chromosphere with an electron number density of cm 3 is about 200 km (Emslie 1978), close to the thickness of the H line center formation layer (Qu & Xu 2002); low-energy protons with an energy lower than 400 kev may be at the origin of the observed radial polarization. However, return-current electrons associated with the propagation of high-energy electron beams could also generate such radial polarization (Hénoux & Karlický 2003) in both lines. The presence of impulsive hard X-ray emission and radio pulses in the event as reported by Hénoux & Karlický (2003) supports the hypothesis of bombardment by high-energy electrons. In order to neutralize the local plasma during the bombardment, return currents are generated. Lower energy return-current electrons are more efficient than beam electrons in exciting the hydrogen lines by collision. They generate polarized line radiation as long as their energy E R is higher than the line excitation threshold E thr. The linear polarization direction depends mainly on the relative value of the return-current electron energy E R compared to the turnover energy E tov (Karlický &Hénoux 2002; Hénoux & Karlický 2003). Charge neutralization along the beam propagation in the solar atmosphere requires the equality ne RV R ¼ ne BV B ¼, where ne B and nr e are the electron number densities in the beam and in the return current, respectively, V B and V R are the velocities of the beam and return-current electrons, and is the particle number flux to be satisfied.

6 Fig. 6. I, P, and spectral profiles in islands where the polarization is radial in H (left) andh (right). In each example, on the left of the profile panels, the distribution of the polarization degree at line center (solid line) along the slit is plotted for comparison with the amplitude of the intensity gradient in arbitrary units (dotted line). The wavelength dependence of the maximum polarization degree (dotted line) observed along the slit is shown, together with the intensity ( full line) and the orientation (dashed line) spectral profiles. The units ( Int, Pol, and Ori) used for the intensity, polarization degree, and orientation angle are given in the top left of each panel. The radial and tangential directions in units of 180 are also plotted (solid line).

7 624 XU ET AL. Vol. 631 Consequently, the energy, E R, of a return electron is given by 8 " E R ¼ 2m e 2 1 #2 9 < V Te n 1=2 e n 2 exp = e : V Te n e ; ; ð5þ where m e is the electron mass. This equation is valid as long as the electric field E does not exceed the Dreicer field. At low electron number densities E exceeds E Dr and we use for continuity V R ¼ 2 1 n e and E R ¼ 2m e 2 1 n 2 : ð6þ e Fig. 7. Examples of the I, P, and spectral profiles in H at locations of tangential polarization at line center. Notations are the same as in Fig. 6. Presumably, only runaway electrons of density n run carry the return current, with n run (Norman & Smith 1978) given by 8 " n run ¼ 1 2 n e exp 1 # 9 < E 1=2 Dr E 2 = : 2 E ; ; ð2þ E Dr where the Dreicer field E Dr is the electric field E that would be present if all background electrons were moving with the thermal velocity V Te.SinceE / j,wherej is the current density, and j / n e V e,then E Dr E ¼ V T e n e : ð3þ Equations (2) and (3) and the particle flux conservation equation lead to 8 " V R ¼ 2 1 exp 1 # 9 < V Te n 1=2 e 2 = n e : 2 V Te n e ; : ð4þ These formulae differ by factors of 2 and 4, respectively, from the expression obtained in the hypothesis that the return current is carried by all background electrons [V R ¼ ð1/n e Þ and E R ¼ 1 2 m e 2 1/n 2 e ]. The differences between these two pairs of formulae reflect the uncertainty of the theory of return-current formation in high electric fields. Whatever the theory applied, the return-current electron energy depends only on the local temperature, the electron number density, and the beam particle number flux. The electron number density dependence of E R for three values of the electron number flux (10 17,10 18,and10 19 cm 2 s 1 ) is plotted in Figure 9. Here E R does not depend on the individual energy of the beam particles. However, in order to relate the number flux values to energy fluxes, knowing that the electron energy flux rarely reaches ergs cm 2 s 1 in solar flares, it is worth giving the energy fluxes of monoenergetic beams made of 10 kev electrons corresponding to the particle number fluxes used. They are equal to 1:6 ; 10 9,1:6; 10 10,and1:6; ergs cm 2 s 1, respectively. The return-current electron energy could be above or below the excitation threshold, and if above this threshold it could be below or above the turnover energy. Consequently, depending on the local conditions and the particle number flux, return-current electrons generate radial or tangential polarization or no polarization at all. To this polarization, the polarization created by the beam electrons must be added. Their effect is dominant in the return-current electron null polarization regions, where they lead to a net weak tangential polarization. The density layers at the boundary between radial and tangential polarization regions are indeed regions of low or no polarization. In an atmosphere in hydrostatic equilibrium, the electron number density increases with depth, and Figure 9 does qualitatively represent the depth dependence of the polarization generated by return current. Consequently, situations may arise in which the layers of formation of H and H line center intensities correspond, respectively, to the layers of null radial or tangential polarization in one or two lines, explaining the origin of the observed polarization islands. Tangential polarization in one line (TN ). Return-current and beam electrons could generate tangential polarization. However, the tangential polarization observed in the TNZ islands is clearly associated with upward motions. This result is in agreement with the observation by Firstova et al. (2002) of H tangential linear polarization in a surge, i.e., orthogonal to the polarization azimuth in the associated flare. Therefore, the origin of the observed tangential polarization must be in some way related to the presence of upward motions. A modified Fletcher & Brown (1998) model could generate some tangential polarization. These authors suggested that hydrogen impact excitation in the regions of interaction of hot evaporating and cool nonevaporating material could generate locally impact-polarized line emission. In their model, due to the difference in vertical upward velocity of the hot and cool

8 No. 1, 2005 IMPACT POLARIZATION OF FLARE 625 Fig. 8. Time evolution of the H and H line intensities two-dimensional spatial distribution at line center and from left to right at k ¼ 0:11 and nm for H and 0.08 and nm for H. These wavelength shifts, interpreted as Doppler shifts, correspond to upward or downward velocities along the line of sight of 50 km s 1. The beginning scan times are 10:07:20, 10:08:54, and 10:10:29 UT from top to bottom, respectively. Positions of the main radial (diamonds) and tangential (asterisks) polarization islands are superposed to compare the distributions of polarization orientation and longitudinal velocity field. The colors of the symbols (white or black) are chosen in order to get the highest contrast on the two-dimensional intensity maps. The H tangential polarization peak appears to be located in an area where a transverse component about 70 km 1 of the upward velocities is present. plasmas, radial polarization was generated. Since the proton energy associated with a 100 km s 1 upflow is only about 50 ev, i.e., much below the 400 kev turnover energy above which the line impact polarization is tangential, such a model built to explain radial polarization could not be used directly to explain a tangential one. If we assume that all cool and hot columns of the evaporating chromosphere are moving up with the same velocity, impact polarization would be due only to interactions Fig. 9. Dependence of the energy E R of a return-current electron on electron number density of F1 flare model ( Machado et al. 1980), for three values of the electron number flux () as given on the figure, in an isothermal atmosphere at temperature T ¼ 10 4 K. T, R, and N are regions of tangential, radial, and null polarization, respectively, generated by return current. perpendicular to the column axis. The temperature gradient at the interfaces of the hot and cool columns would generate tangential polarization. However, in order to get a significant degree of polarization, the radius of each column must be close to the thickness of the interaction region, i.e., below 1 km, requiring a highly filamented chromosphere, and the resulting interpenetration of the hot and cool gas will lead to an enhancement of the H intensity that is not observed. The origin of the observed tangential polarization must be looked for in the mechanisms offormation of solar surges. Electromagnetic j ^ B forces generated in magnetic reconnection have been suggested to be at the origin of surges (Shibata 2001). The tangential direction of polarization observed is by definition perpendicular to the transverse component of the magnetic field, and it could result from hydrogen collisional excitation by the electrons of a current perpendicular to the magnetic field. These electrons carrying this current of density j are presumably taken from the tail of the local electrons thermal velocity distribution. If their densitydecreases strongly with energy, at least in the reported example, they can contribute to H line excitation without having significant effects on the H line. Such a mechanism must be privileged in the explanation of the tangential polarization observed. It is also worth noting, in parallel with differences in the degree of polarization, that, during all the observing time, the longitudinal hydrogen velocities in the whole active region are much larger and last longer in the H line formation layers than in the layers where the H line is formed. In conclusion, the polarization observed can be interpreted as impact excitation due to two sources of collisional anisotropy: (1) particle bombardment by either low-energy protons or highenergy electrons in the lower flaring kernel and (2) in the surge close to the upper flare kernel, the presence of electric currents that generate in the magnetic field the electromagnetic forces that lift the surge.

9 626 XU ET AL. Vol. 631 The authors are pleased to thank Carine Briand of Meudon Observatory for useful codes to treat the basic data. The authors further wish to acknowledge the support of a grant from the French Ministère de la Recherche et de la Technologie and the support of National Key Fundamental Research Project (G and G ) and NSFC ( and ) of China. APPENDIX ALIGNMENT OF THE I þ S AND I S SIGNALS After correcting for line curvature, the exact relative positioning of I þ S and I S spectra was done, together with the correction of a slight difference in the size of the two spectra. Accurate relative positioning of the up and down spectra carrying the I þ S and I S signals, respectively, is mandatory. As shown below, a bad alignment generates a false polarization signal S g proportional to the intensity gradient along the slit direction. At a given wavelength and at a given location i on the Sun, the polarization signal is obtained by subtracting the signal I down ( y i ) along the slit carrying (I S ) from the signal I up ( y þ i )carrying(i þ S ). Here y i and y þ i are the ordinates on the CCD of the same solar Fig. 10. Checking the validity of the value of y. The values of the Stokes parameter U at H line center, obtained for y (middle panel )andy 0:1pixels(bottom and top panels), are shown as a solid line. The gradient-generated Stokes signals Sg 1 for y y ¼1are represented by dashed and dotted lines, respectively. The 1% polarization degree levels are noted by two parallel thin lines. At locations indicated by arrows 1 and 2 where the Stokes signal is 0, a variation from+0.1to 0.1 of the error in relative positioning generates signals with opposite sign, as expected from gradient-generated signals such as Sg 1. Then a significant polarization signal such as the one indicated by arrow 3 does not change significantly for a 0.1 error in positioning.

10 No. 1, 2005 IMPACT POLARIZATION OF FLARE 627 location i seen in the two I S and I þ S spectra, and y (y ¼ y þ i y i ), the separation of these spectra in the y-direction, is generally independent of i. Due to the intensity gradient, even for a null Stokes signal if the two spectra are not relatively well positioned, a false polarization signal S g, as given by equation (2) below, can be generated: S g ¼ I up( y þ i ) I down ( y i ) I up ( y þ i ) þ I down ( y i ) ¼ 1 di (y y): 2I mean dy ða1þ Here I mean is the mean value of I up and I down, di/dy is the intensity gradient, and y is the estimate of the spatial displacement y. For a given error on the relative spectra positioning, (y y), S g is proportional to (1/I )(di/dy). Thebestestimateofyis the one that makes the polarization in the H or H line wings the smallest, after all necessary corrections are done. The intensity gradient is particularly high when the slit is crossing small sunspots, and this positioning is the most appropriate location to compare the polarization signal to the intensity gradient. The false signal due to the intensity gradient there gives useful hints on how to remove the positioning error and to find the right value of y for Q and U. In Figure 10, the variation with position y along the slit of the polarization signal at H line center is used in order to check the validity of the y estimate. For this estimate of y (middle panel)andy 0:1pixels(bottom and top panels), the y-dependence of the Stokes parameter U (solid line) is compared with the gradient polarization Sg 1 due to y y ¼1(dashed line/dotted line). At locations indicated by arrows 1 and 2, a significant intensity gradient is present; nevertheless, a null Stokes signal is found there when using the exact y value. Then errors in relative positioning as small as +0.1 to 0.1 generate signals with opposite signs, as expected from gradient-generated signals such as Sg 1. Indeed, as indicated by arrow 3, a Stokes signal much larger than the gradient polarization does not change significantly for a 0.1 error in positioning. Aboudarham, J., Berrington, K., Callaway, J., Feautrier, N., Henoux, J. C., Peach, G., & Saraph, H. 1992, A&A, 262, 302 Bommier, V., & Rayrole, J. 2002, A&A, 381, 227 Emslie, A. G. 1978, ApJ, 224, 241 Firstova, N. M., Hénoux, J. C., Kazantsev, S. A., & Bulatov, A. V. 1997, Sol. Phys., 171, 123 Firstova, N. M., & Kashapova, L. K. 2002, A&A, 388, L17 Firstova, N. M., Xu, Z., & Fang, C. 2003, ApJ, 595, L131 Fletcher, L., & Brown, J. C. 1995, A&A, 294, , A&A, 338, 737 Hanaoka, Y. 2003, ApJ, 596, 1347 Hénoux, J.-C., Chambe, G., Smith, D., Tamres, D., Feautrier, N., Rovira, M., & Sahal-Brechot, S. 1990, ApJS, 73, 303 Hénoux, J.-C., & Karlický, M. 2003, A&A, 407, 1103 Hénoux, J.-C., & Semel, M. 1981, in Proc. Solar Maximum Year Workshop, ed. V. N. Obridko & E. V. Ivanov ( Moscow: Izmiran), 207 Karlický, M., & Hénoux, J.-C. 2002, A&A, 383, 713 REFERENCES Kašparová, J. 2004, Ph.D. thesis, Charles Univ., Czech Republic Kazantsev, S. A., & Hénoux, J.-C. 1995, Polarization Spectroscopy of Ionized Gases ( Dordrecht: Kluwer) Kleinpoppen, H. 1969, in Advances in Atomics and Molecular Physics, ed. F. Bopp & H. Kleinpoppen (Amsterdam: North-Holland), 612 Machado, M. E., Avrett, E. H., Vernazza, J. E., & Noyes, R. W. 1980, ApJ, 242, 336 Norman, C., & Smith, R. 1978, A&A, 68, 145 Qu, Z. Q., & Xu, Z., 2002, Chinese J. Astron. Astrophys., 2, 71 Shibata, K. 2001, Encyclopedia of Astronomy and Astrophysics ( Bristol: IOP) Vogt, E., & Hénoux, J.-C., 1996, Sol. Phys., 164, , A&A, 349, 283 Vogt, E., Sahal-Bréchot, S., & Bommier, V. 2001, A&A, 374, 1127 Werner, A., & Schartner, K.-H. 1996, J. Phys. B, 29, 125 Xu, Z., Firstova, N. M., Chen, Q.-R., & Fang, C. 2003, Chinese J. Astron. Astrophys., 3, 266