ANALYSIS OF A SOLAR CORONAL BRIGHT POINT EXTREME ULTRAVIOLET SPECTRUM FROM THE EUNIS SOUNDING ROCKET INSTRUMENT

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1 The Astrophysical Journal, 677:781Y789, 2008 April 10 # The American Astronomical Society. All rights reserved. Printed in U.S.A. A ANALYSIS OF A SOLAR CORONAL BRIGHT POINT EXTREME ULTRAVIOLET SPECTRUM FROM THE EUNIS SOUNDING ROCKET INSTRUMENT Jeffrey W. Brosius Catholic University of America at NASA Goddard Space Flight Center, Solar Physics Laboratory, Code 671, Greenbelt, MD 20771; jeffery.w.brosius@nasa.gov Douglas M. Rabin and Roger J. Thomas NASA Goddard Space Flight Center, Solar Physics Laboratory, Code 671, Greenbelt, MD 20771; douglas.m.rabin@nasa.gov, roger.j.thomas@nasa.gov and Enrico Landi Artep, Inc., Naval Research Laboratory, Code 7660, Washington, DC 20375; enrico.landi@nrl.navy.mil Received 2007 September 12; accepted 2007 December 19 ABSTRACT We present a well-calibrated EUV spectrum of a solar coronal bright point observed with the Extreme Ultraviolet Normal Incidence Spectrograph ( EUNIS) sounding rocket instrument on 2006 April 12. The coronal bright point brightened around 06:30 UT during a period of emerging magnetic flux and remained bright at least until the rocket flight around 18:12 UT, while the magnetic flux merged and canceled. Density-sensitive line intensity ratios yield mutually consistent coronal electron densities (N e in cm 3 )oflogn e 9:4. The differential emission measure (DEM, in cm 5 K 1 ) curve derived from the spectrum yields a peak of log DEM 20:70 at log T 6:15 and a local minimum of log DEM 20:15 at log T 5:35. Photospheric (not coronal) element abundances are required to achieve equality and consistency in the DEM derived from lines of Mg v,mg vi,mg vii,and Ca vii (with a low first ionization potential, or FIP) and lines from Ne iv and Ne v (with a high FIP) formed at transition region temperatures. The bright point s photospheric abundance is likely produced by reconnection-driven chromospheric evaporation, a process that is not only central to existing bright point models, but also consistent with measurements of relative Doppler velocities. Subject headinggs: Sun: corona Sun: magnetic fields Sun: transition region Sun: UV radiation Online material: color figuresonline MATERIALCOLOR FIGURESFIGURE 1FIGURE 3 1. INTRODUCTION Coronal bright points were discovered in soft X-ray images of the solar corona obtained with sounding rocket instruments (Vaiana et al.1970). They were first studied in detail with observations from the S-054 X-ray spectrographic telescope aboard Skylab, which revealed mean lifetimes of about 8 hr, a typical maximum area of 2 ; 10 8 km 2, and X-ray filter ratio temperatures between 1.3 and 1:7 ; 10 6 K (Golub et al.1974). Based on observations acquired by the Harvard EUV spectrometer/spectroheliometer, Habbal et al. (1990) found that fewer bright points reached coronal temperatures than reached transition region temperatures. More recently Zhang et al. (2001) investigated bright points using simultaneous images obtained with the Yohkoh satellite s Soft X-Ray Telescope and the Solar and Heliospheric Observatory (SOHO) spacecraft s Extreme-Ultraviolet Imaging Telescope (EIT; Delaboudinière et al.1995) and found smaller numbers of bright points ( 1 3 to 1 4 as many) in soft X-rays (for which the plasma temperature exceeded 2:5 ; 10 6 K) than in EUVemission of Fe xii at (for which the plasma temperature is about 1:4 ; 10 6 K). The average lifetime of the bright points that were observed from initial appearance (in Fe xii emission) to final disappearance was 20 hr. McIntosh & Gurman (2005) developed an automated detection algorithm to derive properties of bright points observed by SOHO s EIT during nine years of operation and found nearly 100 times more bright points in EIT s Fe ix/x wave band (T 1 ; 10 6 K) than in its Fe xv wave band (T 2 ; 10 6 K). Thus, hotter bright points (like the one discussed in the present work) are much less common than cooler ones. 781 While all magnetic bipoles are not necessarily associated with coronal bright points, all coronal bright points are associated with photospheric magnetic bipoles ( Krieger et al. 1971; Golub et al. 1977). This association suggests that bright points are produced by coronal magnetic reconnection initiated when opposite-polarity photospheric magnetic fragments continuously approach one another ( Priest et al. 1994; Parnell et al. 1994a, 1994b; Longcope 1998; Longcope et al. 2001; von Rekowski et al. 2006a, 2006b). The resulting energy released into the corona leads directly to the formation of a bright point by local heating and/or chromospheric evaporation. This scenario further suggests that bright points are dynamic phenomena, in which relative Doppler velocities are expected to be observed. Several investigators have used SOHO s SUMER instrument (Solar Ultraviolet Measurements of Emitted Radiation; Wilhelm et al. 1995) to measure Doppler velocities in bright points. ( Here and in what follows we write T m as the temperature in K that maximizes the relevant ion s fractional abundance.) For example, Vilhu et al. (2002) measured upflow speeds of 20 5kms 1 from blueshifted emission of N iv at (T m 1 ; 10 5 K); Madjarska et al. (2003) measured Doppler velocities from 10 to +10 km s 1 across a bright point with the S vi line at (T m 2 ; 10 5 K); Popescu et al. (2004) measured velocities between 4 and +5 km s 1 in a bright point with the Mg ix line at (T m 1 ; 10 6 K). More recently, Brosius et al. (2007) used coronal bright point spectra obtained with the Extreme Ultraviolet Normal Incidence Spectrograph ( EUNIS) sounding rocket instrument to measure relative upflow and downflow velocities of 26 km s 1 in Fe xiv at (T m 2:0 ; 10 6 K) and

2 782 BROSIUS ET AL. Vol km s 1 in Fe xvi at both and (T m 2:5 ; 10 6 K), the hottest lines for which Doppler velocities have been reported in a bright point. Velocities in He ii at (formed at T m 5 ; 10 4 K) and Mg ix were found to be 15 and 14 km s 1, respectively. Here we carry our investigation of the coronal bright point observed with EUNIS further by presenting and analyzing an averaged EUV spectrum from the core of the bright point. Lines in the spectrum enable thorough temperature coverage for 5:2 P log T m P 6:4. We use the integrated line intensities to derive coronal electron densities, the differential emission measure ( DEM), and elemental abundances in the bright point. To the best of our knowledge this is the first time that the DEM and elemental abundances have been presented for a coronal bright point. In x 2we describe the observations and data reduction procedures, in x 3we provide the results of our analysis, in x 4 we discuss the implications of our findings, and in x 5wesummarizeourconclusions. 2. OBSERVATIONS AND DATA REDUCTION We observed a coronal bright point with the EUNIS sounding rocket instrument around 18:12 UT on 2006 April 12 and with SOHO s EIT and Michelson Doppler Imager ( MDI; Scherrer et al.1995) throughout the day until shortly after the rocket flight. The EIT and MDI observations reveal that EUV brightenings in the bright point were associated with emerging magnetic flux and that the bright point s emission remained elevated as the magnetic flux subsequently diminished. EUNIS observed the bright point during this phase of elevated EUV emission, roughly 11 hr after the start of the long-duration brightening. Based on a visual inspection of full-disk images obtained in all four of EIT s wave bands during and around the time of the rocket flight, the coronal bright point investigated here appears to be rather ordinary: at any given wavelength, it does not stand out in either size or brightness compared to other bright points in the same images EIT Observations We obtained a full-disk He ii image with SOHO s EITat 18:12 UT (near the start of the EUNIS observations) and 18:20 UT (shortly after the end of the EUNIS observations). The EUNIS lobe images of He ii at , obtained from the wide areas observed at the ends of the narrow slit (see Brosius et al. 2007, and below), were co-aligned with the EIT He ii image at 18:12 UT to determine the EUNIS slit position and drift during the stare mode. See Figure 1, where the nearly horizontal solid lines indicate the position of the EUNIS slit at the start (bottom line) and the end (top line) of the stare sequence. Based on different methods that we used for co-alignment, we estimate an uncertainty P3 00 on the slit s y-position in this figure. We used the 12 minute cadence EIT Fe xii images to quantify the evolution of the bright point (maximum intensity, area within a specified intensity threshold, and average intensity within that area) from the beginning of April 12 until shortly after the rocket flight. These images reveal that the bright point both brightened and enlarged significantly during the course of the day. See Figure 2. Because we do not display the bright point s area in this figure, we mention for reference that it started the day with an area of 1:09 ; 10 8 km 2, decreased to a minimum of 5:82 ; 10 7 km 2 around 6:00 UT, and increased to 5:02 ; 10 8 km 2 by 18:06:47 UT (shortly before the EUNIS flight) MDI Observations We used the full-disk photospheric longitudinal magnetograms obtained with MDI at a cadence of 96 minutes to investigate the evolution of the bright point s magnetic field. (Although MDI Fig. 1. Coronal bright point image in He ii at obtained with EIT at (a)18:12 UT (near the start of the EUNIS stare observations) and (b)18:20 UT (shortly after the rocket flight), where x- and y-axes show solar disk coordinates in arcseconds (as viewed from 1 AU). The lower of the two nearly horizontal lines in each frame shows the pointing of the EUNIS long-wave slit at the start of the sequence of 75 1 s stare exposures; the upper one shows the slit at the end of that sequence. Contour levels indicate count rate increases by factors of 2. Some variability in the bright point morphology is evident, but no flaring or drastic intensity changes occurred. [See the electronic edition of the Journal for a color version of this figure.] actually obtained a magnetogram every minute throughout the day, an analysis of the bright point s field variation on such short timescales is beyond the scope of the present investigation: here we seek mainly to place the EUNIS observations within the context of the overall evolution of the bright point.) Each 96 minute magnetogram was obtained simultaneously with one of the Fe xii images acquired with EITat a 12 minute cadence. To determine the photospheric longitudinal magnetic fields within the bright point pixels identified in the EIT images, we rebinned the magnetograms so that their pixel size matched that of the EIT images and co-aligned the resulting full-disk magnetograms with their corresponding Fe xii images. Then, for the same bright point pixels identified in the EIT images, we obtained (1) the maximum value of the outward-directed magnetic field strength, (2) the average value of the outward-directed field (for bright point pixels within which the field was directed outward), (3) the minimum value of the inward-directed magnetic field strength, and (4) the average value of the inward-directed field (for bright point pixels within which the field was directed inward). The top frame of Figure 2 shows the maximum absolute values of the inward-directed (dashed curve) and outward-directed (dotted curve) photospheric longitudinal magnetic field strengths (in G) measured with MDI, along with the maximum value of the bright point s Fe xii emission (solid curve) measuredwith EIT (in counts s 1 ) as a function of time (hours) since 00:00 UT on 2006 April 12; the bottom frame shows the corresponding average values. The bright point clearly began to brighten (and expand in area) between 06:00 and 07:00 UT, with the emergence of new negative magnetic flux. The bright point s Fe xii emission remained elevated until the EUNIS flight even though its magnetic field strength ( both positive and negative) was steadily declining

3 No. 1, 2008 SOLAR CORONAL BRIGHT POINT EUV SPECTRUM 783 Fig. 2. Evolution of the coronal bright point s Fe xii emission observed with EIT (in counts s 1, shown as a solid curve), and its outward-directed (dotted curve) and absolute value of its inward-directed (dashed curve) photospheric longitudinal magnetic field strength (in gauss) obtained with MDI from the beginning of the day until just after the EUNIS flight on 2006 April 12. Frame (a) displays maximum values of the above three quantities observed in the bright point, and frame (b) displays their average values; for the inward-directed and outward-directed average field strengths, values were calculated using only those bright point pixels for which the field was appropriately directed. The dot-dashed line indicates the time of the EUNIS flight. (likely due to merging and cancellation). Thus, our bright point initially brightened during a phase of emerging magnetic flux, and sustained itself as that flux subsequently merged and canceled EUNIS Observations A brief description of the EUNIS sounding rocket instrument is provided by Brosius et al. (2007). We mention here only those characteristics that are pertinent to the present investigation. The EUNIS long-wave channel covers first-order wavelengths between 300 and 370 8, with a throughput of nearly 100 times that of its SERTS (Solar EUV Research Telescope and Spectrograph; Neupert et al. 1992; Thomas & Neupert1994; Brosius et al. 1996, 1998, 2000a, 2000b) predecessor. Count rates are recorded with three 1024 ; 1024 active pixel sensors (Stirbl et al. 1998), such that small gaps in wavelength coverage occurred around 321 and The EUNIS spectral dispersion is 25 m8 pixel 1, and its spectral resolution, although designed to achieve 100 m8 full width at half-maximum (FWHM; comparable to that of SERTS) was 200 m8 FWHM on this first flight. The spectrograph s entrance aperture is shaped roughly like an hourglass, with a narrow 1:71 00 ; slit along which spatially resolved spectra are obtained, and wider lobes at the ends of the slit from which spectroheliograms are imaged. The latter are used to precisely co-align the EUNIS data with full-disk He ii images obtained with EIT. The EUNIS spatial pixel size is , but because the optics limited its actual spatial resolution to about 5 00 on this first flight, we averaged over three spatial pixels for a net spatial pixel size of EUNIS was launched from White Sands Missile Range, New Mexico, at 18:10 UT on 2006 April 12 and acquired solar spectra and images between 18:12 and 18:18 UT. During the first half of its flight EUNIS obtained 75 1 s stare exposures (fixed pointing) with a cadence of 2.10 s (several exposure intervals were somewhat longer than this owing to computer overhead). Although the slit remained tilted about 1.5 counterclockwise from solar disk x-axis during the flight, its pointing drifted about per exposure in both the x- and y-directions, for a net drift of about in each direction during stare mode. Figure 1 shows the pointing of the EUNIS slit during the first (the bottom nearly horizontal solid line) and the 75th (the top nearly horizontal solid line) 1 s stare exposure atop a portion of the EIT He ii image obtained at 18:12 UT (Fig. 1a) and 18:20 UT (Fig.1b), displayed on a negative intensity scale. The net drift during stare mode was sufficient to carry the slit almost off the bright point. Therefore, for the present work, we average over only the first 15 stare exposures to obtain the bright point core spectrum. We processed the EUNIS spectral data by subtracting an average of dark frames obtained in flight, debiasing each detector row with a linear fit to the unilluminated ends, and applying a flatfield image obtained in the laboratory. We applied the absolute radiometric calibration derived from postflight end-to-end measurements carried out at Rutherford Appleton Laboratory ( RAL), and incorporated corrections for atmospheric extinction. The uncertainty on the absolute radiometric calibration is 10%. As with most solar spectrometers, EUNIS does not provide an absolute wavelength scale. However, we constructed a relative wavelength scale based on the known spectral dispersion and selected reference lines. We derived integrated intensities for lines observed in the EUNIS spectra using standard Gaussian line-profile fitting procedures developed for analysis of data from SOHO s Coronal Diagnostic Spectrometer (CDS; Harrison et al. 1995), available

4 784 BROSIUS ET AL. Fig. 3. Integrated line intensities of (a)heii at ,(b)Mgix at , and (c)fexiv at obtained with EUNIS in and around the coronal bright point, displayed as a space-time plot. The x-axis shows solar disk coordinates in arcseconds (as viewed from 1 AU), and the y-axis shows time (in seconds) since the start of the first 1 s stare exposure. Rectangular boxes indicate the 5 spatial by 15 temporal pixels from which the averaged bright point spectrum displayed and analyzed in this work was obtained. [See the electronic edition of the Journal for a color version of this figure.] through IDL s SolarSoftware. For each wavelength bin that we selected for emission-line intensity integration, profile fits were made to one or more Gaussians on a linear background, resulting in values of the centroid position, integrated intensity, and FWHM for each line, together with their corresponding statistical uncertainties. An additional 10% systematic uncertainty associated with the absolute radiometric calibration is also included in the uncertainty on the integrated intensity. Figure 3 shows a space-time image of integrated line intensities of He ii at , Mgix at , and Fe xiv at obtained with EUNIS in and around the bright point, displayed on a negative intensity scale. The x-axis corresponds to position along the slit, and the y-axis corresponds to time, with the first 1 s stare exposure along the bottom of each frame and the 75th along the top. Boxes in each frame outline the space-time portion of the EUNIS spectral data from which we obtained the averaged bright point spectrum shown in Figure 4 and described in x RESULTS We present the first (to the best of our knowledge) EUV spectral line list for a solar bright point. From it we derive not only electron densities but also the differential emission measure (DEM) and elemental abundances for the bright point Coronal Bright Point EUV Spectrum Figure 4 shows the coronal bright point spectrum obtained by averaging the EUNIS spectral data from the five ( ) spatial pixels and the 15 exposures outlined by the rectangular boxes in Figure 3. We consider this limited spatial and temporal subset of the EUNIS observations to comprise the core of the bright point. Dotted vertical tick marks indicate emission lines, whose integrated intensities are listed in Table 1. The hottest lines in this table are those of Fe xvi, which suggests an upper limit 2:5 ; 10 6 K for the bright point s temperature. We used the bright point spectrum to verify the EUNIS relative radiometric calibration by comparing observed values with theoretical values of line intensity ratios that are insensitive (or only slightly sensitive) to density and temperature ( insensitive ratios ). A number of insensitive ratios are available for the lines of Ne iv,nev,mgv,mgvii,mgviii,siviii,fexi,fexii,fexiii, and Fe xvi listed in Table 1. These are useful not only to verify the calibration, but also to confirm line identifications and blends in the spectrum. Here we mention only a few such ratios between lines that are relatively far apart in the EUNIS spectral range. We derived theoretical values for the ratios from the CHIANTI (ver. 5.2; Dere et al.1997; Landi et al. 2006) database. The ratios vary very slowly (if at all) with density, but for completeness we estimated the uncertainty on the theoretical ratio as half the difference between the maximum and minimum values provided by CHIANTI for densities between 10 8 and cm 3. For example, for the intensity ratio of the two Mg viii lines at and we find that the observed value of 0:24 0:04 is in excellent agreement with the theoretical value of 0:23 0:02. Similarly, for the lines of Fe xii at and the observed and theoretical ratios of 0:37 0:06 and 0:31 0:01 agree within their uncertainties. The lines of Fe xi at and yield observed and theoretical ratios of 0:30 0:05 and 0.30, while the lines of Fe xvi at and yield observed and theoretical values of 0:45 0:06 and The agreement between these and additional observed and theoretical insensitive ratios demonstrates the reliability of the EUNIS relative radiometric calibration. We also verified the EUNIS absolute radiometric calibration by comparing our He ii plus Si xi quiet-sun intensities against previous measurements. To do this we used Fe xiv and He ii space-time images similar to those in Figure 3 to identify two separate slit segments in the quiet Sun, one to the east of ( but separated from) the bright point and one to the west of (also separated from) the bright point. Within this area, we find an average He ii plus Si xi blended line intensity of ergs cm 2 s 1 sr 1, where the uncertainty is based on the 10% associated with the EUNIS laboratory calibration. This intensity is consistent with the average quiet-sun intensity of ergs cm 2 s 1 sr 1 for this He ii plus Si xi line pair adopted by Mango et al. (1978) based on well-calibrated Orbiting Solar Observatory series satellite and sounding rocket observations, and very close to the disk center value of 7115 ergs cm 2 s 1 sr 1 that they derive in their limb-brightening study using the Skylab slitless spectrograph. Our quiet-sun He ii plus Si xi intensity of ergs cm 2 s 1 sr 1 is smaller than (even accounting for uncertainties), but comparable to, the He ii ( ) plus Si xi

5 TABLE 1 Average Intensities in the Core of the Coronal Bright Point Observed by EUNIS Line log T m FWHM Intensity Fig. 4. Average spectrum of the coronal bright point observed with EUNIS, separated into wave bands corresponding to each of the active pixel sensors. Frame (a) displays the spectral intensity in ergs cm 2 s 1 sr on a logarithmic scale because the He ii line at is very intense. Frames (b) and(c) display the spectral intensity in ergs cm 2 s 1 sr on a linear scale. Dotted vertical tick marks indicate all of the lines (except Fe xiii at ) whose integrated intensities are listed in Table 1. Si xi He ii Fe xi Fe xiii Mg viii Si viii Mg viii Si viii Mg viii Fe xiii Mg vii Si viii Fe xiii Unid Unid Fe xv Unid Cr xiii Fe ix Al x Fe xiv Mg viii Fe xvi Fe xii Mg viii Fe xi Si ix Ca vii Unid Fe xii Si x Fe xiii Mg vi Si ix Mg v Fe xii Fe xi Mg v bl Fe xiv Si x Fe xi Ne iv Ne v Ne iv Ne v Fe xiii Fe xiii Fe xvi Unid Fe xii Mg vii Ne v Mg vii Mg ix Fe xi Notes. Wavelengths and full width at half-maximum ( FWHM) values are in 8. T m is the temperature at which a given line s fractional ion abundance is maximum, here adopted from Mazzotta et al. (1998) as incorporated into the CHIANTI (ver. 5.2) database (Dere et al. 1997; Landi et al. 2006). Line intensities are in ergs cm 2 s 1 sr 1.TheMgvi line is a blend of four closely spaced Mg vi lines. All of the unidentified lines have been observed previously with SERTS-89 and/or SERTS-97 and/or CDS.

6 786 BROSIUS ET AL. Vol. 677 (643 84) intensity reported by Brosius et al. (2000b) in the quiet-surroundings spectrum observed by SERTS in 1997 from the vicinity of an active region. The version of SERTS flown in 1997 (SERTS-97) was also absolutely radiometrically calibrated at RAL. We used the CHIANTI database to determine the theoretical density dependence of density-sensitive line intensity ratios. Then, using the observed ratios derived from line intensities in Table 1, we find the following coronal electron densities in the bright point. For Fe xi, the ratio of intensities of the line to the line (i.e., Fe xi /352.67) yields log N e ¼ 9:42 þ0:13 0:15.ForFexii / we obtain log N e ¼ 9:69 þ0:21 0:27,forFexiii / we obtain log N e ¼ 9:54 þ0:12 0:11,forFexiv / we obtain log N e ¼ 9:52 þ0:13 0:15,andforSix / we obtain log N e ¼ 9:40 þ0:28 0:23. All of these values (corresponding to lines formed at 6:1 P log T m P 6:3) are mutually consistent. The coolest density-sensitive ratio available in our data is that of Mg vii / (log T m ¼ 5:8), which yields log N e ¼ 9:00 þ0:13 0:13.Ugarte- Urra et al. (2005) derived bright point electron densities from line intensity ratios for some of these same ions observed in six bright points with SOHO s CDS. For Fe xiv, for example, their average measured density in the six bright points extends from 9:20 þ0:33 9:50 þ0:04 0:04 0:83 to 0:13 to 9:26þ0:15 0:14.,whileforSix it ranges from 8:97þ0:13 They concluded that their measurements suggested that bright point plasma is more similar to active region than to quiet-sun plasma The Differential Emission Measure The observed intensity I ij of an optically thin line emitted by an electron transition from upper level j to lower level i of the ion X þm can be written I ij ¼ 1 Z 4 T GT; ð N e Þ ðtþdt photons cm 2 s 1 sr 1 ; where ðtþ is the differential emission measure of the plasma, defined as ðtþ¼ Ne 2 dh dt cm 5 K 1 ; ð1þ ð2þ and G(T; N e ) is the contribution function of the emitting line, defined as GT; ð N e Þ¼ N jðx þm Þ NðX þm Þ NðXÞ NðX þm Þ NðXÞ NðHÞ NðHÞ A ji cm 3 s 1 ; ð3þ N e N e where N j (X þm )/N(X þm ) is the relative population of the upper level j, N(X þm )/N(X) is the relative abundance of the ion X þm, N(X)/N(H) is the abundance of the element X relative to H, N(H)/N e is 0.83 for fully ionized plasmas of standard composition, and A ji is the Einstein coefficient for spontaneous emission from level j to level i. The contribution function includes all the atomic physics involved in the process of line formation, and the DEM provides information about the temperature distribution of the plasma in the emitting source. The determination of the DEM from a set of optically thin lines requires the inversion of equation (1) for all the observed lines in the data set, once the contribution function of each of them has been calculated as a function of temperature. In order to avoid uncertainties due to the density sensitivity of G(T; N e ), density-insensitive lines need to be used in DEM studies, or N e needs to be measured before determining the DEM. Several techniques have been developed to invert equation (1) and determine the DEM using different methods and approximations. Here we use the iterative technique developed by Landi & Landini (1997), in which the DEM is written ðtþ¼ 0 ðtþ! ðtþ; where 0 (T) is an initial estimate of the DEM curve and!(t )is a correction function assumed to vary smoothly with temperature. Landi & Landini (1997) demonstrated that for a given set of line intensities their final DEM solution is independent of their initial estimate. In the present work we use a quiet-sun DEM (unpublished, but similar to the quiet-sun DEM curves of Brosius et al. 1996) for 0 ðtþ. The inversion of equation (1) for each line k allows us to calculate the value of! k ðt ea Þ at an effective temperature ð4þ R T log T ea ¼ GT; ð N eþ 0 ðtþlog TdT R T GT; N : ð5þ ð eþ 0 ðtþdt The corrections! k (T ea ) provided by all the lines in the data set are first averaged in temperature bins of appropriate size [(log T bin ) ¼ 0:3] to smooth out any irregularity given by lines with problems due to unknown blends or inaccurate atomic physics, then interpolated over temperature to determine!(t ). The result is used in equation (4) to calculate a new 1 (T ) trial function to be used in the next iteration. This technique provides the final DEM curve when the function!(t ) is unity within uncertainties. In order to apply the Landi & Landini (1997) technique to the line intensities observed with EUNIS, we first calculated the contribution functions using version 5.2 of the CHIANTI database, adopting the ion fractions of Mazzotta et al. (1998). We repeated the calculation using both the photospheric elemental abundances of Grevesse & Sauval (1998) and the coronal abundances of Feldman et al. (1992), but reserve a discussion of our final selection of the most appropriate set of abundances for x 3.3. The lines observed by EUNIS are mostly coronal, plus a few transition region lines from Ne iv, Ne v, Mg v, Mg vi, Mg vii, and Ca vii. Many of the coronal lines, especially those from Fe xiyxiv and Si x, are density sensitive, so that in order to minimize density effects in the DEM determination we used an electron density consistent with the values derived above in our calculations (log N e ¼ 9:4). Some of the lines in Table 1 were excluded from our DEM analysis either because they were so buried in the wings of much stronger nearby neighbors that their intensities could not be reliably determined (even though their listed formal uncertainties do not look so bad), or because they consisted of blends of lines formed at very different temperatures. In the first case, we discarded Si xi and Mg viii because they were in the wings of the much stronger He ii and Fe xvi lines. (Historically, this is one of the reasons for treating the sum of the He ii and Si xi intensities in assessments of absolute calibration as described above.) In the second case, we rejected the lines of Fe xii+mg v at , Neiv+Fe xi at , and Ne v+fe x at because their calculated log T ea is found at an intermediate temperature between the log T m values of the two ions contributing to the blend, so that the values of! k (T ea ) provided by these lines influence the DEM at temperatures far from those at which their measured intensity is actually formed. The two Fe xiii lines at and were treated as a single line, and their measured intensities were summed together to minimize the uncertainties associated with separating their profile fits.

7 No. 1, 2008 SOLAR CORONAL BRIGHT POINT EUV SPECTRUM 787 TABLE 2 Coronal Bright Point DEM log T (K) log DEM (cm 5 K 1 ) Fig. 5. Coronal bright point DEM derived from the EUNIS spectral line intensities in Table 1, based on photospheric abundances. Values derived from lines of different elements are displayed with different symbols as shown in the figure. Note, in particular, the agreement among transition region lines of low-fip Mg and high-fip Ne. The lines that lie far below the curve around log T ¼ 6:0 are those of Si viii (at , , and ), which ion has also revealed similar problems in spectra from SOHO s SUMER. We adopted a standard value of (log T bin ) ¼ 0:3 to calculate the DEM displayed in Figure 5. This curve shows the typical features of a rather quiet DEM with a broad peak log DEM 20:70 centered around log T 6:15,andaminimumlogDEM 20:15 at log T 5:35. There are a few lines whose intensities are not well reproduced by the DEM in Figure 5, the most notable of which are the three Si viii lines around log T ¼ 6:0, whose emissivities are too large by a factor 2.4. These lines are strongly density sensitive, but they all depend on N e in the same way so they cannot be used to determine N e. A possible cause of this disagreement might be problems in the ion fractions of Si viii: Landi et al. (2002) reported an anomalous behavior of N-like ions relative to all other ions and ascribed that to ion fractions; similar problems were found by Landi & Feldman (2008) in active regions. Alternatively, it is possible that the DEM is oversmoothed so that Si viii could be formed in a narrow temperature region where the DEM has a deep minimum. To test this we decreased the width of (log T bin ) to 0.05 in order to increase the temperature resolution of the DEM. The resulting curve shows two narrow peaks at log T ¼ 5:8 and 6.1, and a minimum at log T ¼ 5:95 that brings predicted Si viii line intensities in agreement with the rest of the observations. However, even though the 2 of the solution is slightly lower than with (logt bin ) ¼ 0:3, the solution is unstable because it does not converge to a single final solution (it oscillates between two or three different curves), and therefore it cannot be accepted. In order to actually achieve a finer T bin resolution we need a larger number of ions in the data set to allow for a finer sampling of the temperature range. Lines from Si ix are also overestimated relative to Fe xi lines at around log T ea 6:1 by a factor 2.1. This ion, when coupled with Si viii, suggests that Si abundance might be overestimated by a factor 2Y2.5, but the agreement of Si x lines with those from other ions formed at similar temperature leads us to think that Si viii and Si ix ion fractions are overestimated while the adopted relative Si/ Fe element abundance is correct. The DEM curve (Fig. 5) at log T < 5:2 and log T > 6:4 is not constrained by our data set because EUNIS observed lines formed only in the range 5:2 log T m 6:4(whichcorrespondstoNeiv through Fe xvi; see Table 1). This means that the DEM outside the EUNIS temperature range is nearly unaltered from its initial Note. This table lists log DEM values only for the temperature range actually covered by emission lines observed with EUNIS. estimate, except that it must be connected smoothly and continuously to the constrained portion of the curve. We do not use the He ii line at , formed at log T m ¼ 4:7, in our DEM analysis, because that line is not only optically thick, but may also be affected by photoionization/recombination and/or velocity redistribution in addition to collisional excitations ( Zirin 1988; Jordan et al. 1993; Andretta et al. 2000; Jordan & Brosius 2007). In Table 2 we list DEM values for 5:2 log T 6:4, i.e., only those that are actually derived from the well-calibrated EUNIS spectrum of the coronal bright point Element Abundances As described above, we calculated the bright point s DEM from the observed EUNIS line intensities for each of two different sets of elemental abundances: the photospheric abundances of Grevesse & Sauval (1998) and the coronal abundances of Feldman et al. (1992). The difference between these two sets of abundances depends on the first ionization potential ( FIP) of the various elements. The FIP effect ( Feldman & Laming 2000, and references therein) is a phenomenon in which coronal abundances of low-fip ions (<10eV;e.g.,Na,Mg,Al,Si,Ca,Cr,Fe, Ni) are enhanced (by typical factors 4) relative to their photospheric values, while coronal abundances of high-fip ions (>10 ev; e.g., He, C, N, O, Ne) are the same as in the photosphere. All of the coronal lines available in the present EUNIS data set belong to low-fip ions, but the available transition region lines include a mixture of low-fip ( Mg, Ca) and high-fip ( Ne) ions. This combination, therefore, can be used to study the FIP effect in the bright point. The set of abundances for which the derived DEM is self-consistent ( particularly among lines of different ions formed at and near the same temperature) is the one that most likely corresponds to reality. Of particular interest here is the fact that Mg v and Ne v have nearly the same log T m (see Table 1) and log T ea (5.5), which means that these two different ions should yield the same value for the DEM (within uncertainties) at this log T ea.otherlow-fipions with log T m (see Table 1) and log T ea close to that of Mg v and Ne v include Ca vii and Mg vi, while other high-fip ions include Ne iv. Even though we cannot directly compare values for the DEM derived from these ions (because they are not at exactly the same log T ea ), we can use the resulting smoothness and continuity of the derived DEM to assess the quality of the curve

8 788 BROSIUS ET AL. Vol. 677 derived from the two different sets of abundances. We find, first, that results based on the Mg v and Ne v lines agree only when photospheric abundances are used; they do not agree with each other (by a factor of about 4) when coronal abundances are used. We find, further, that the DEM at transition region temperatures is much smoother and reproduces the Ca and Mg line intensities much better with the Grevesse & Sauval (1998) photospheric abundances than with the Feldman et al. (1992) coronal abundances. Thus, we conclude that the element abundances in the coronal bright point observed with EUNIS were photospheric. 4. DISCUSSION Our conclusion regarding photospheric element abundances in the coronal bright point observed with EUNIS provides new observational evidence that supports reconnection models of bright points ( Priest et al. 1994; Longcope1998). In these models, coronal magnetic reconnection is initiated when opposite-polarity photospheric magnetic fragments continuously approach one another ( Priest et al. 1994; Parnell et al. 1994a, 1994b; Longcope 1998; Longcope et al. 2001; von Rekowski et al. 2006a, 2006b). This releases energy into the corona that provides both local heating and acceleration of particle beams. Chromospheric evaporation occurs when the beamed particles (or thermal conduction front or both) deposit energy into the chromosphere more rapidly than it can be radiated away, thus causing the chromosphere to heat and expand. The resulting evaporation can be either explosive or gentle ( Fisher et al.1985a,1985b,1985c; Brosius & Phillips 2004) depending on the rate of energy deposition; the two can be distinguished by their respective temperature-dependent velocity patterns. Either way, chromospheric evaporation dredges up material from the chromosphere or lower in the atmosphere, where the FIP bias is not observed to occur. Therefore, our observation of photospheric element abundances in the coronal bright point is consistent with expectations that chromospheric evaporation should produce just such an effect. Measurements of Doppler velocities associated with chromospheric evaporation observed in EUV lines with SOHO s CDS are given by Brosius & Phillips (2004) for a solar flare and by Brosius & Holman (2007) for a flarelike transient. Brosius & Phillips (2004) measured upward velocities 40 km s 1 in lines formed at transition region temperatures during gentle evaporation and downward velocities 40 km s 1 in those same lines during explosive evaporation (when hot flare lines formed around 10 7 K revealed simultaneous upward velocities 70 km s 1 ). Brosius & Holman (2007) measured simultaneous, cospatial downward velocities 30 km s 1 in the chromospheric line of He i at and the transition region line of O v at , and upward velocities 20 km s 1 in the coronal line of Si xii at during explosive chromospheric evaporation. These values are comparable to the upward and downward velocities of 35 km s 1 (for Fe xvi), 26 km s 1 (Fe xiv), 14 km s 1 (Mgix), and 15 km s 1 (Heii) reported by Brosius et al. (2007) across the bright point observed by EUNIS. Thus, our measured velocities across the bright point are consistent with previous values associated with chromospheric evaporation. Our conclusion regarding photospheric element abundances in the coronal bright point observed with EUNIS is also consistent with abundance evolution observations reported by Feldman & Widing (2007). They argue that the FIP bias in the solar upper atmosphere changes at a slow but steady rate based, in part, on spectroscopic observations of six active regions that were monitored by Widing & Feldman (1995) from the time they were born until they decayed or rotated beyond the west limb. Soon after birth the abundances in the active regions were photospheric, i.e., they had a FIP bias of 1. After 2Y3 days the FIP bias values reached 4Y5, and after about 5 days the values reached 8Y9. Feldman & Widing (2007) point out that in some old active regions and in some flaring plasmas the FIP bias can reach values as high as 16; however, in the average quiet-sun corona and in the slow-speed solar wind the FIP bias is typically 4Y5. It is interesting to note that White et al. (2000), who used coordinated EUV and radio observations of an aged active region in which the coronal magnetic field strength was insufficient to produce thermal gyroemission, derived a coronal iron abundance 4Y5 times that of photospheric. Since we observed not only the emergence of new magnetic flux but also the EUV brightening of the bright point earlier (by about 11 hr) during the day of the EUNIS flight, we might have anticipated that our bright point s abundances would be photospheric. The bright point DEM derived in this work reveals a peak of log DEM 20:70 at log T 6:15 and a local minimum of log DEM 20:15 at log T 5:35. By way of comparison, the active region DEM derived from the well-calibrated spectrum obtained with SERTS-97 (Brosius et al. 2000a) reveals log DEM 21:35 (a factor of 4 greater than that of the coronal bright point) at log T 6:15, and log DEM 21:00 (a factor of 7 greater than that of the bright point) at log T 5:35. Further, the SERTS-97 active region DEM shows a peak of log DEM 21:60 at log T 6:30 and a local minimum of log DEM 20:50 at log T 5:75. At these same log T values, the bright point DEM yields log DEM 20:35 and 20.30, which are factors of 18 and 1.6, respectively, less than the corresponding active region values. Thus, the active region DEM exceeds that of the coronal bright point at all temperatures over which the two curves are well defined, but the difference is smallest at transition region temperatures near the minimum of the active region s curve. 5. SUMMARY We have presented a well-calibrated EUV spectrum of a solar coronal bright point observed with the EUNIS sounding rocket instrument. SOHO observations reveal that the bright point brightened during a period of emerging magnetic flux and remained bright for at least 11 hr (until the rocket flight), while the magnetic flux merged and canceled. Density-sensitive line intensity ratios yield mutually consistent coronal electron densities log N e 9:4 in lines formed at 6:1 P log T m P 6:3. The DEM curve derived from the spectrum yields a peak of log DEM 20:70 at log T 6:15 and a local minimum of log DEM 20:15 at log T 5:35. We find that photospheric (not coronal) element abundances are required to achieve equality and consistency in the DEM derived from transition region lines of low-fip Mg v, Mg vi,mgvii,cavii, and high-fip Ne iv and Ne v. The coronal bright point s photospheric abundance is likely produced by reconnection-driven chromospheric evaporation, a process that is not only central to existing bright point models, but also consistent with measurements of relative Doppler velocities. The EUNIS program is supported by the NASA Heliophysics Division through its Low Cost Access to Space Program in Solar and Heliospheric Physics. We thank the entire EUNIS team for the concerted effort that led to a successful first flight. The work of E. Landi is supported by NNG06EA14I, NNH06CD24C, and other NASA grants. A careful reading and valuable comments by the referee are appreciated.

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