A high temperature corona above an active region complex

A HIGH TEMPERATURE REGION COMPLEX

CORONA ABOVE AN ACTIVE

Y.-K. Ko i, J. C. Raymond i, J. Li 2, A. Ciaravella i,3, J. Michels 1,4, S. Fineschi i'5, and R. Wu i

iHarvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138, USA 2Institute for Astronomy, University of Hawaii, 2680 Woodlawn Dr., Honolulu, HI 96822, USA 3now at Osservatorio Astronomico di Palermo, "G.S. Vaiana", P.za Parlamento 1, 9013~ Palermo, Italy 4now at Princeton Materials Institute, Princeton University, Princeton, NJ 085~, USA 5now at Osservatorio Astronomico di Torino, Strada Osservatorio 20, 1-10025, Pino Torinese, Italy

ABSTRACT We present the results of SOHO/UVCS and Yohkoh/SXT observations above an active region complex (AR8194/8195/8198) at the southeast limb on April 6-7, 1998. The electron temperature analysis indicates a two-temperature structure, one with ~ 1.5 • 106 K which is similar to that observed in quiet Sun streamers, the other with a high temperature ~ 3.0 • 106 K. We compare the electron temperature and emission measure from the SOHO/UVCS data with those from the Yohkoh/SXT data. The absolute elemental abundances show a general first ionization potential effect (FIP effect) and decrease with height for all the elements. We discuss mechanisms that may explain the observed abundances.

OBSERVATION AND RESULTS The target of the observation is the corona above an active region complex at the southeast limb. AR8194 and AR8195 appeared at the east limb on April 5 followed by AR8198 1.5 days later. The heliographic latitude is S18, $27 and $28 for AR 8194, 8195, and 8198, respectively. UVCS observed for nine heliocentric heights centered at position angle (PA) of 120 ~ from 13:27 UT, April 6 to 14:20 UT, April 7, 1998. Figure 1 shows the lowest and highest OVI Channel slit positions (1.22 R| to 1.60 R+) on the EIT-UVCS composite image, and the Yohkoh/SXT image plotted with the regions extracted for SXT data analysis. UVCS data show that besides the usual coronal lines, some high-ionization lines, such as Fe XVII )~1153, [Fe XVIII] A974, Ne IX )~1248, [Ca XIV] A943, are particularly bright compared to the quiet Sun corona. These high-ionization lines are still visible at heights up to 1.6 R+. This indicates that this region is unusually hot compared with the 'average' solar corona observed at these heights. The UVCS slit is 40 arcmin (.-~ 2.5R| long as seen in the plane of the sky (see Figure 1). We average UVCS data over 280 arcsec of the spatial extent of the slit centered at PA=120 ~ (corresponding to PA range from ~ 114 ~ to ~ 126~ This is where the emission of those high-ionization lines is concentrated. The data have been wavelength and radiometrically calibrated, and corrected for stray light and flat field. The radiative and collisional components for the hydrogen Lyman lines and OVI A1032/A1037 doublets are separated. The collisional excitation rates are mostly adopted from the CHIANTI database version 3.01 (Dere et al., 2001). The ionization equilibrium of Mazzotta et al. (1998) were adopted. For details of the UVCS data analysis, see Ko et al. (2002). -73-

Y.-K. Ko et al.

Fig. i. Left panel: The pointing of UVCS observations on the composite image of EIT 284 (19:06 UT, April 6) and UVCS (synoptic image in OVl ~I032, taken from 22:06 UT, April 5 to 11:50 UT, April 6, 1998). Right panel: Yohkoh/SXTimage at 23:36 UT, April 6 plotted with the regions extracted for data analysis. Figure 2 plots the electron temperature derived from various line ratios of Si and Fe lines at all heights. It can be seen that the temperature distribution seems mainly to be clustered into two temperatures, one around 1.5 x 106 K, and the other around 3 x 106 K. If we assume that the corona above this active region complex has a two-temperature structure along the line of sight, and that the elemental abundances in the two Te regions are the same, the line intensity (photon s-lcm-2sr -1) can then be expressed as:

l nel nion / Izine = --47-~HBline[ne---7(Thi)qline(Thi)( nenHdl)hi + ni~176176

/

nenHdl)lo]

(1)

where net~nil is the elemental abundance relative to hydrogen (absolute abundance), nion/net is the ionic fraction, Bline is the branching ratio, qline is the electron excitation rate, and f nenHdl is the emission measure at a given Te. Thi and ~o are the average of the ratio temperatures (see Figure 2) from the high-Te gas (FeXV/FeXVII, FeXVIII/FeXVII, FeXVIII/FeXV) and the low-Te gas (rest of the ratios), respectively. If we use [Fe X] A1028, [Fe XII] A1242 as the proxy for the low-Te gas and Fe XVII Al153, [Fe XVIII] A974 for the high-Te gas, the ratio of the emission measures at the two Te's can be determined. The 'high-Te' and the 'low-Te' components for the lines can then be calculated analytically using Eq. 1 along with the absolute elemental abundance and the emission measure. Figure 3 plots the abundances relative to their photospheric values versus their first ionization potentials (FIP). The lines that are used to determine the adopted abundances are: N V A1238, O VI A1032, Ne IX A1248, Si XII A499, [S X] )~1196, [Ar XII] A1018, [K XIII] A994, [Ca XIV] A944, the average of [Fe X] A1028, [Fe XII] A1242, [Fe XIII] Abl0, Fe XVII Al153 and [Fe XVIII] A974 for Fe, and [Ni XIV] A1034. We can see that the FIP effect, in which the abundances of the low-FIP elements (FIP smaller than ,,~10 eV, such as Fe) are enhanced relative to those of the high-FIP elements (FIP larger than ,-~10 eV, such as Ar) when compared with their photospheric values, is present at all heights. Furthermore, the abundance generally decreases with height in a systematic way. -74-

A High Temperature Corona above an Active Region Complex I

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Figure 4 and Figure 5 compare the electron temperature and the emission measure, respectively, derived from the UVCS data with those from the SXT data which are taken from 09:09 UT, April 6 to 11:09 UT, April 7, 1998. The data were averaged over PA=120 4- 5~ with increment of 0.05 Re (cp. Figure 1). The SXT Science Composite (SSC) images with filter pairs AI.1 and A1Mg were used. In order to compare the SXT and UVCS measurements, we calculate theoretical X-ray spectra developed by Raymond & Smith (1977) ('RS code') using the elemental abundances measured by UVCS ('ab_UVCS'). The SXT response function is obtained by combining the theoretical X-ray spectra with the SXT effective area. The SXT temperature and emission measure are then derived from the SXT response function. The comparison shows that the electron temperatures of the 'high-Te' region from UVCS are consistent with those from SXT. The temperature derived from SXT band ratios is expected to be an average of the high and low temperatures, strongly weighted toward the high temperature by the higher emissivity of the hotter plasma in the Yohkoh bandpasses. The emission measures derived from SXT are also consistent with the UVCS results. Figures 4 and 5 also show the SXT results using two other approaches: 1) SXT standard routine ('SXT code') which uses Meyer (1985) coronal abundances ('ab_Meyer'), and 2) Raymond & Smith (1977) code using Meyer (1985) coronal abundances. We can see that it is important to use consistent elemental abundances and plasma emission codes to calculate the electron temperature and the emission measure from broad band data such as SXT acquires. Using different sets of abundances (e.g. photospheric vs. coronal, coronal with FIP effect with enhanced vs. depleted low FIP elements) may give significantly different results.

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Y.-K. Ko et al. '

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Fig. 3. Elemental abundances relative to the photospheric values vs. their first ionization potential at all nine heights.

DISCUSSION AND SUMMARY We have shown that the corona above this active region complex can be characterized by electron temperature higher than that in the usual quiet Sun corona and streamers. Detailed analysis indicates that the temperature distribution is mainly clumped around two values, one at ~ 1.5 x 106 K and the other at ,~ 3 x 106 K. The lower value is similar to that found in the quiescent streamers. The higher value is most likely associated with these active regions. Our analysis shows that the FIP effect is present at all heights and the abundance decreases with height for all the elements. The FIP bias of about a factor of 4 is typical of the slow solar wind. Schwadron et al. (1999) modeled the elemental fractionation at the foot of large coronal loops and found that MHD wave heating is able to provide both mass-independent fractionation and low-FIP bias in coronal loops. The materials stored in the closed loops are then released by reconnection with adjacent open field lines and form the slow wind carrying FIP bias with them. A plausible explanation for the decreasing abundances at larger heights is gravitational settling of the heavier elements in closed magnetic loops (Raymond et al. 1997). We have shown that the absolute abundance declines by factors of 2-4 between 1.2 and 1.6 Ro, suggesting a scale height of about a few 10 l~ cm. There is no obvious dependence upon mass. Gravitational settling, however, cannot be totally responsible for the abundance variations we see here. Active regions usually evolve on a time scale of a few weeks. The settling times of the ions are roughly one day (Lenz et al., 1998). Therefore -76-

A High Temperature Corona above an Active Region Complex 6.80 6.70

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Fig. 4. The electron temperature derived from UVCS compared with those from SXT using three different approaches.

we would expect much smaller abundances than we observed here. One possibility to counteract the effect of gravitational settling is that plasma may cycle through the magnetic loops with a time scale of a few days. Another possibility is that the apparently static gas is flowing outwards, and the heavy ions are pulled along by ion drag of the protons (e.g. Ofman 2000). This latter option would suggest that very highly ionized plasma would be observable by ACE at times corresponding to the passage of these active regions. The decline in abundance with height would also suggest an increase with height in the ratio of outflow speed of the elements to the speed of hydrogen. Elemental abundance is a powerful tool in understanding the coronal origin of the solar wind. This could be accomplished by comparing the elemental abundances in the corona and those in the solar wind measured in-situ. Previous and present work have shown that FIP effect (which is based on relative abundances of low-FIP to high-FIP elements) exists in the streamers and active region loops. However, the absolute abundances change with height and across structures. Therefore, absolute abundances, not the FIP effect solely, should be the more relevant parameter in understanding the coronal origin of the solar wind. A CKN OWLED G EMENT S This work is supported by NASA grant NAG5-10093.

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K-K. Ko et al.

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REFERENCES Dere, K. P., Landi, E., Young, P. R., and Del Zanna, G., Astrophys. J. Suppl., 134, 331 (2001) Ko, Y.-K., Raymond. J. C., Li, J., Ciaravella, A., Michels, J. et al., Astrophys. J., in press (2002) Lenz, D. D., Lou, Y.-Q., and Rosner R., Astrophys. J., 504, 1020 (1998) Mazzotta, P., Mazzitelli, G., Colafrancesco, S., and Vittorio, N., Astron. Astrophys. Suppl., 133, 403 (1998) Meyer, J.-P., Astrophys. J. Suppl., 57, 173 (1985) Ofman, L., Geophys. Res. Lett., 27, 2885 (2000) Raymond, J. C., et al., Sol. Phys, 175, 645 (1997) Raymond, J. C., & Smith, B. W., Astrophys. J. Suppl., 35, 419 (1977) Schwadron, N. A., Fisk, L. A. and Zurbuchen, T. H., Astrophys. J., 521, 859 (1999)

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