upper transition region lines observed close to the solar limb to be emitted from UFSs and to show quasiphotospheric

Transcription

1 THE ASTROPHYSICAL JOURNAL, 518:926È936, 1999 June 20 ( The American Astronomical Society. All rights reserved. Printed in U.S.A. THE OFF-LIMB BEHAVIOR OF THE FIRST IONIZATION POTENTIAL EFFECT IN T [ 5 ] 105 K SOLAR PLASMAS J. M. LAMING,1 U. FELDMAN,2 J. J. DRAKE,3 AND P. LEMAIRE4 Received 1998 June 16; accepted 1999 January 22 ABSTRACT We investigate the behavior of the solar Ðrst ionization potential (FIP) e ect (the abundance enhancement of elements with Ðrst ionization potential of less than 10 ev in the corona with respect to photospheric values) with height above the limb in a region of di use quiet corona observed by the SUMER instrument on SOHO, with emphasis on so-called upper transition region lines. Previous disk observations have shown di erent abundance patterns in emission from lines at temperatures above and below D8 ] 105 K, with an FIP e ect clearly visible at the higher temperatures and less so at the lower temperatures. Our initial aim is to determine whether such a di erence is also visible in o -limb observations. We Ðnd a low-fip element enhancement of a factor of 3È4 indicated in all line ratios. The Mg VII /Ne VII ratio is also seen to decrease toward a photospheric value when tracked down from the corona to the limb. This is markedly di erent from the behavior of higher temperature line ratios and may be related to the di ering heating and mass supply mechanisms for plasmas at temperatures above and below D8 ] 105 K. An additional unexpected feature of our observations is that in low-fip/high-fip line ratios formed at temperatures close to the freeze-in temperature of the fast solar wind (D106 K), there is also a small diminution of the FIP enhancement at the highest altitudes observed in this study. We discuss the possible relevance of this to the origin of the fast solar wind. Subject headings: Sun: abundances È Sun: corona È Sun: transition region È Sun: UV radiation 1. INTRODUCTION Despite extensive investigation over the past three decades or so, the structure of the solar upper transition region and corona (hereafter referred to as the solar upper atmosphere) remain enigmatic. In active regions, coronal plasma (with electron temperature T [ 8 ] 105 K) is e known to reside in loops that must be connected to the lower lying and much colder chromosphere. The same is generally believed to be true elsewhere in quieter regions, though not in coronal holes. The thin region of temperature intermediate between chromospheric and coronal values has become known as the transition region and traditionally has been interpreted as the source of solar emission from ions formed at these characteristic temperatures. Although such a thin transition region must exist in coronal loops, Feldman & Laming (1994) have argued that the observed emission from these regions is masked by other magnetic structures, named the unresolved Ðne structures ÏÏ (UFSs) by Feldman (1983, 1987), which contain plasma with temperatures \8 ] 105 K. Alternative scenarios for this thin transition region have been explored by WikstÔl et al. (1998). The element abundances in coronal loops and in the slow-speed solar wind are by now fairly well established to be subject to the FIP (Ðrst ionization potential) e ect, (see e.g., Feldman 1992) whereby elements with Ðrst ionization potential of less than 10 ev have their abundances 1 SFA, Inc., Largo, MD E. O. Hulburt Center for Space Research, Code 7608, Naval Research Laboratory, Washington, DC Smithsonian Astrophysical Laboratory MS 70, 60 Garden Street, Cambridge, MA Institut dïastrophysique Spatiale, Unite Mixte CNRS-Universite Paris XI-Bat 121, F91405 Orsay Cedex, France. 926 enhanced by a factor 3.5È4 in the corona with respect to photospheric values. The fact the low-fip elements are enhanced, and that high- FIP elements are not depleted relative to hydrogen, has been established from the O/H ratio by Feldman et al. (1998). A recent investigation of the solar upper transition region and coronal abundances derived from full-disk solar observations revealed a hitherto unnoticed feature of the FIP e ect. Laming, Drake, & Widing (1995), working from the spectrum of Malinovsky & Heroux (1973), observed the usual FIP e ect in true coronal plasma with temperatures º106 K, but at lower temperatures emission from the upper transition region plasma indicated photospheric abundances. Similar e ects are visible, but not commented on, in data analyzed by other workers (Noci et al. 1988; Doschek, Dere, & Lund 1991). This photospheric abundance material was assumed to reside in the unresolved Ðne structures, largely because of the correspondence in the temperatures for these structures deduced by Feldman (1983, 1987) and those where photospheric abundances were observed by Laming et al. (1995). The Solar and Heliospheric Observatory (SOHO) o ers a new opportunity to investigate further the relationship between the putative unresolved Ðne structures and coronal loops, by studying the height dependence of element abundances and the FIP e ect o -limb. Implicit in this strategy is the assumption that the UFSs exist only at low altitudes, as is suggested, e.g., by observations with SO82A (see Feldman & Widing 1990, 1993; Widing & Feldman 1989, 1992, 1993), and the Harvard EUV spectrometerspectroheliometer (Vernazza & Reeves 1978). We expect upper transition region lines observed close to the solar limb to be emitted from UFSs and to show quasiphotospheric abundances, giving way to coronal abundance emission from coronal loops at higher altitudes. True coronal lines should show coronal abundances at all posi-

2 OFF-LIMB BEHAVIOR OF THE FIP EFFECT 927 tions, inside, on, and outside the solar limb, since the UFSs do not contribute to any of this emission. We establish Ðrst the degree of the FIP e ect in lines typically emitted from T \ 8 ] 105 K plasmas high up in the corona. Then, if this shows e a pattern di erent from that previously established for the disk emission, we try to image the transition from one abundance pattern to the other as one moves from the corona to the limb. In this way we can hope to glean extra insight into the FIP e ect and its rami- Ðcations for the heating and mass supply for the corona and solar wind by trying novel observations with the new facilities a orded by SOHO. 2. OBSERVATIONS The data used here were taken with the SUMER instrument on SOHO (Wilhelm et al. 1995, 1997; Lemaire et al. 1997) on 1997 May 5È6. The 0A.3] 120@@ slit was used, cen- tered at (]699A.9, ]699A.9) in a north-south orientation. An image of the Sun taken on 1997 May 5 at 22:12 UT in the light of predominantly Fe XII 195 A by the EIT instrument on SOHO (Delaboudinière et al. 1995) is shown in Figure 1, with the position of the SUMER slit superposed. The slit axis was aligned in the north-south direction, making an angle of 45 to the solar limb. Each spatial pixel along the slit corresponds to approximately 1/J2 arcsec (D500 km radial distance) above the solar limb. The slit intercepted about 20A inside the limb and extended to about 100A above it. The region imaged was di use quiet corona, well away from speciðc solar regions like coronal holes or active regions. The main difficulty in observing line ratios from the socalled transition region is the time variability of this emission. The standard SUMER reference spectra are taken with an integration time of 5 minutes per detector position, FIG. 1.ÈSOHO EIT image taken on 1997 May UT illustrating the position of the SUMER 120A slit for the observations discussed herein. Notice that the lower end of the slit is inside the solar limb.

3 928 LAMING ET AL. Vol. 518 TABLE 1 LINE SELECTIONS Central Wavelength Target Lines FIP (A ) (A ) log T (ev) Ne VIII , Mg VIII , , Al VIII Mg IX S X Ne VII , Na VII Mg VII S IX Si IX AlVII , Si VII S VII Ar XII Ne VI (]2) Al XI (]2) SX , Mg VI , Mg VII Si XI SiVIII , this being the time taken to telemeter the data from one exposure back to the receiving station. Within this 5 minute time, substantial variability (approximately a factor of 2 or more) can occur in transition region lines formed at temperatures up to D3 ] 105 K (see Feldman et al. 1997), with fewer indications of variability at higher temperatures. However, according to the UFS picture, we expect lines formed at temperatures up to D8 ] 105 K to be emitted from similar structures, so we searched for detector positions that would allow us to record lines from low- and high-fip elements formed as close as possible in temperature within the D43 A Ðrst order bandpass of the SUMER detector. Table 1 gives the central wavelengths of these detector positions. In practice not all such low-fip/ high-fip intensity ratios could be observed simultaneously, so observational sequences were repeated to obtain a time average. These selected detector positions also allowed us to observe suitable density diagnostic line ratios to monitor other plasma conditions. These were the Mg VI / , Ne VII /895.17, Al VI / , Si VIII / , and S X / density diagnostic ratios. Figures 2aÈ2d show the G(T ) contribution functions (the product of the ionization balance in collisional ionization equilibrium, taken from Arnaud & RothenÑug (1985), and upper level excitation rate as a function of temperature, taken to be P exp ([*E/kT )/T 1@2) for various of the transition region ions observed in this study. Note the particularly good correspondence between Mg VII , Al VII , Si VII , and S VII Similar correspondences are found between the various ions for which the temperature of maximum fractional abundance is in the coronal regime T [ 8 ] 105 K. However, in the coronal e temperature region the emission measure distribution FIG. 2.ÈTheoretical G(T ) contribution functions for the transition region and lower coronal ions used in this study. The following transitions are plotted: (top left)mgvi A and Ne VI A ; (top right)nevii A, Mg VII A, Al VII A, Si VII A, and S VII A ; (bottom left) Ne VIII A, Mg VIII A, and Al VIII A ;(bottom right)mgix A, S IX A, Si IX A, and S X A.

4 No. 2, 1999 OFF-LIMB BEHAVIOR OF THE FIP EFFECT 929 changes much less rapidly with temperature than in the lower temperature region, and hence the precise overlap of G(T ) functions is less critical. 3. DATA ANALYSIS The detector image Ðles were Ñat-Ðelded and destretched ÏÏ to correct for small-scale detector sensitivity variations and image distortions using standard SUMER software routines. Strong unblended lines can be measured easily and ratioed by simply extracting the appropriate region of spectrum, subtracting the background, and summing over the line proðle. In this way, abundance and density diagnostic ratios shown below are determined as a function of distance above the solar limb. Weaker and/or blended lines require Ðtting model line proðles to the data in order to determine their intensities Electron Density Diagnostics Figures 3a, and 3b show the N-like density diagnostic line ratios Si VIII and S X (see Laming et al for a full discussion of this and other isoelectronic sequences). In each case, line intensities (in units of counts per spatial pixel) integrated over the line proðle and over the 1 hr exposure) and their ratios are plotted as a function of distance along the slit in the upper panel. The lines are distinguished by dashes (Si VIII and S X ), dash-dots (Si VIII and S X ), solid lines for the ratios, and dotted lines for the error bars on the intensity ratios. The lower panel shows detector images of the lines aligned so that the direct spatial correspondence between the image and the plot may be seen. The solar limb is at the right of each image, visible as a band of bright continuum. All images are aligned with respect to the southernmost end FIG.3a FIG. 3.ÈDetector images and plots of density diagnostic line ratios for (a) SiVIII A / A and (b) SX A / A. Each image is aligned with the plot above it, so that the spatial pixel number can be read for both the plot and the image. The solar limb is at the right hand side, about 25È30 pixels in from the edge. The pixel numbers vary slightly, since the image of the 120@@ ] 0A.3 falls on slightly di erent positions on the detector for di ering wavelengths. On the plots the line intensities are indicated by dashed and dash-dot lines, the former corresponding to the upper image in the lower panel (i.e., , or A ), and the dash-dot line corresponding to the lower image. Each image is 40 pixels (D1.68 A ) wide, with wavelength increasing from top to bottom. The ratios in each case, plotted as a solid line (with 1 p statistical uncertainties as dotted lines), can be combined with theoretical curves to infer an logarithmic electron densities of 8.15 ^ 0.1 and 8.5 ^ 0.1, respectively, for (a) and (b). At each end of the plot spurious intensities and intensity ratios can result. At the left-hand end we reach the end of the exposed part of the detector, and on the right-hand side the higher continuum and contributions from colder lines emitted from the disk can make background subtraction very difficult. FIG.3b

5 930 LAMING ET AL. Vol. 518 FIG.3c FIG.3d of the slit, which is inside the solar limb. Thus exactly the same area of the Sun is imaged in each observation. Each plot and image is the sum of two 30 minute integrations taken during a forward and backward scan through the various detector positions given in Table 1. The intensities plotted are not corrected for spectrometer absolute sensitivity, i.e., they are raw counts taken directly from the detector images shown in the lower panels. The error bars plotted for the ratio (dotted lines) are taken from the Poisson noise in these raw counts and are the 1 p uncertainties for each spatial pixel. In more quantitative work to follow, we sum over regions of 20 spatial pixels, resulting in a considerable improvement in signal-to-noise ratio. The intensity ratios plotted are of course corrected for the relative sensitivities of the SUMER instrument at the relevant wavelengths. To guide the eye, the intensities and ratios have been smoothed over 9 pixels, but the error bars have no smoothing. With reference to Figure 5 in Laming et al. (1997), the Si VIII lines indicate a logarithmic density of 8.15 ^ 0.1 cm~3 at all distances above the limb considered in this paper (i.e., pixels 145È225 in Fig. 3a, top panel). The uncertainty is estimated from the scatter in the observed ratio. This density decreases slightly as one moves to yet higher positions above the solar limb. The S X ratio indicates a higher density; 8.5 ^ 0.1 cm~3, with less indication of a decrease at higher altitudes. The Si VIII densities are consistent with those found by Laming et al. (1997), whereas those for S X are slightly, but probably not signiðcantly, higher. Further electron densities inferred from other line ratios are given in Table 2, measured by line proðle Ðtting to one-dimensional spectra extracted from the May 5È6 detector images after summing over a number of spatial pixels. The precise ranges are R, pixels 1È20; R, pixels 21È40; R, pixels 41È60; and R 1, pixels 61È80, where 2 the pixel numbering 3 goes from right 4 to left on the images, i.e., R covers the solar limb, and R are successively D10,000 1 km further away from it. Close 2,3,4 to the limb it is often very difficult to measure coronal forbidden lines against the chromospheric continuum, but farther away measurements become possible. The values inferred for the densities are consistent with those coming directly from the detector images described above, and their values giving the electron density generally in the range 108 to 109 cm~3 conðrm that we are observing true coronal emission and not just scattered light from the solar disk, where densities of about an order of magnitude higher would be found (see Vernazza & Mason 1978). Some of the lines from which we deduce rela-

6 No. 2, 1999 OFF-LIMB BEHAVIOR OF THE FIP EFFECT 931 TABLE 2 INFERRED OFF-LIMB LOGARITHMIC ELECTRON DENSITIES (IN cm~3) Line ratio (A ) R 1 1 R 2 R 3 R 4 Mg VI / ^ 0.32 \8.1 Ne VII / [ ^ ^ 0.2 Al VII / ^ ^ ^ 0.5 Si VIII / ^ ^ ^ ^ 0.1 S X / ^ ^ ^ ^ R refer to di erent spatial positions from which the intensity ratios were taken. R 1,2,3,4 1 is on the solar limb. R are successively 20 pixels further up in the corona measured along 2,3,4 the slit, which corresponds to about 14A in radial distance from the solar limb. 2 Uncertainties in densities come from the statistical errors in the intensity ratio measurements propagated through the density diagnostic curves. tive element abundances are mildly density sensitive (e.g., Mg VI), and so we also use these densities measured o -limb in the treatment of these transitions. Some restrictions on the use of diagnostic line ratios are discussed by Judge, Hubeny, & Brown (1997). These have mainly to do with the ill-conditioned nature of the derivation of density structure from intensity ratios. We feel that the issues raised by these authors are unlikely to a ect our conclusions, because of the broad agreement among our various density diagnostic line ratios (and in a similar solar region studied by Laming et al. 1997), and the lack of obvious structure in the solar region we observe. In any case, precise knowledge of the density and its structure is unnecessary for our purposes Relative Element Abundances Figures 4a and 4b show similar plots/images to those described above, but for line pairs that may be expected to be sensitive to relative element abundances. We show the intensities and the intensity ratios for the Mg VII A /Ne VII A and the Mg VIII A /Ne VIII A pairs. The second ratio is formed from lines emitted by hotter coronal ions, and as can be seen, this intensity ratio shows no change across the solar limb. In contrast, the Ðrst ratio that is formed by the colder lines does change in moving from the solar disk to above the limb. The Mg VII A line becomes more intense relative to the Ne VII A as one moves higher up in the atmosphere. This behavior is consistent with, but does not prove, the hypothesis discussed in the introduction to this paper. This is the idea that upper transition region lines (T \ 8 ] 105 K) from low-fip elements should show a smaller e abundance with respect to high-fip elements when observed on the solar disk as opposed to o -limb observations. To be more quantitative in our discussion, we give in Table 3 the measured intensity ratios for various upper transition region and coronal abundance diagnostic line ratios for the positions R described in the previous 1,2,3,4 subsection. These ratios indicate coronal abundances (i.e., a FIP e ect of 3 or more) in every position observed. Most of these intensity ratios cannot be measured on or near the solar limb, since the line blending becomes too severe. Even the ratio that hinted at the e ect we are looking for, Mg VII/Ne VII on closer inspection shows closer to coronal than photospheric abundances, even though Mg appears to become even further enhanced in abundance relative to Ne at higher altitudes. In determining the theoretical intensity ratios for photospheric and coronal abundances we have followed the formalism in Laming, Drake, & Widing (1995) whereby the emission measure distribution is assumed to be Ñat over the TABLE 3 INFERRED OFF-LIMB ABUNDANCE RATIOS Observed1 Theory2 Line ratio R 1 1 R 2 R 3 R 4 R p 2 R c 2 Mg VI /Ne VI ^ 8% ^ 6.5% 0.92 ^ 15% Mg VII /Ne VII ^ 10% 0.31 ^ 4% 0.59 ^ 1.8% 0.67 ^ 2% Na VII /Ne VII ^ 25% ^ 6% ^ 15% Al VII /Ne VII ^ 17% 0.15 ^ 2.4% 0.15 ^ 2.6% 0.20 ^ 3.5% Mg VIII /Ne VIII ^ 4.4% ^ 2.1% ^ 2.5% ^ 2.5% Al VIII /Ne VIII ^ 27% ^ 12% ^ 2.3% ^ 9% Mg IX /S IX ^ 22% 5.1 ^ 11% 4.62 ^ 1.6% 4.5 ^ 1.8% Mg IX /S X ^ 5% 2.4 ^ 2.5% 1.95 ^ 1.8% 2.23 ^ 1.8% Si IX /S IX ^ 50% 4.57 ^ 13% 3.35 ^ 1.9% 3.47 ^ 1.9% Si IX /S X ^ 20% 2.13 ^ 8% 1.42 ^ 2% 1.72 ^ 2% Si XI /S X ^ 11% ^ 6% ^ 5% Si XI /Ar XII ^ 40% 14.4 ^ 10% 13.4 ^ 9% Spatial positions are the same as for Table 2. 2 The theoretical ratios are for assumed photospheric and coronal abundances, R and R, respectively, taken from Feldman et al. (1992). p c 3 Percentage errors in observed intensity ratios are derived from counting statistics only (coming from the line Ðtting program).

7 932 LAMING ET AL. Vol. 518 FIG.4a FIG.4b FIG. 4.ÈDetector images and plots of line ratios sensitive to abundance ratios: (a) MgVII A /Ne VII A and (b) MgVIII A /Ne VIII A. The correspondence between images, plots, and line styles is exactly the same as for Fig. 3. G(T ) functions of the ions involved. Hence emission measures are calculated according to the formula EM \ 4n d 2 I/f, (1) _ where d is the distance to the Sun in cm, I is the measured Ñux in photons _ cm~2 s~1 in a line and f is the number of photons s~1 emitted by unit A EM, given by f \ (n j A ji ) ion /2 \ n j A n n 0.8 ji ion A n2 n & n n n 2 e e j j Bion, (2) A H where n and n are the number densities of a particular ionic excited j level e j and of electrons, respectively, and A is the radiative decay rate. The factors n, n, and n are ji the number densities of the ion in question, ion the A element H from which this ion is formed, and of H, respectively. The factor of 0.8 is the ratio of H atoms or ions to electrons, assuming an essentially fully ionized plasma. In considering full-disk spectra, it is usual also to include a factor of 1/2, accounting for photons emitted back toward the solar disk and absorbed by the photosphere. For o -limb spectra this factor will be increased somewhat, but since we are interested only in ratios of emission measures in various lines (to derive abundance ratios), we omit this factor in the following. In a two-level atom approximation we can write f \ An i C j bf(t )A 0.8, (3) & n max el j j Bion where C is the excitation rate from level i to level j, F(T ) is the ionization ij fraction at the peak of the emissivity curve, max and A is the elemental abundance relative to H. The average el shape of the emission function curves is accounted for through the factor b, formally deðned by b \ / G(T )d(log 10 T ), (4) 0.3G(T ) max

8 No. 2, 1999 OFF-LIMB BEHAVIOR OF THE FIP EFFECT 933 FIG.4c FIG.4d where G(T ) is the product of F(T ) and the factor (n C /& n ). For most of the transitions we consider, we solve i ij the j j equations of statistical equilibrium for a model ion, including usually at least 10 levels for the appropriate electron density, which should be more reliable than the simple two-level atom approximation. This Ñat emission measure ÏÏ assumption is probably not adequate in o -limb spectra of upper transition region ions, since the emission measure here is likely to be a steeply rising function of the temperature. Thus line ratios where one ion has a G(T ) function extending to higher temperatures will be weighted more in favor of that ion. Referring to Figure 1, we can see that this is indeed the case for Mg VI and Mg VII, compared to Ne VI and Ne VII, where intensity ratios suggesting unusually large degrees of FIP fractionation are found. We have computed correction factors for the theoretical intensity ratios in Table 3, calculated for various assumed power law behaviors for the emission measure (EM P T a). The canonical ÏÏ slope is a \ 1.5, corresponding to an emission measure dominated by a constant thermal conductive Ñux downward. Such a slope is commonly observed in full-disk spectra (see Laming et al. 1995). Outside the solar limb one might expect the slope to be greater, since relative to the disk spectra there will be a deðcit of cold material. The correction factors at the larger values of a become suspect since these values are particularly sensitive to the behavior of the G(T ) functions well away from their maxima and, consequently, in the temperature regions where they are least accurately determined. To estimate the relevant values of a for the solar regions imaged by our slit, we calculate the intensity ratios for Ne VI, VII, VIII, MgVI, VII, VIII, and Si VII, VIII, IX for di erent power-law emission measures and compare with our observations. In doing this we have to assume that the element abundance with respect to hydrogen does not change with position. Thus the Ne ions are the most important series, since Ne, being high FIP, does not change in abundance relative to hydrogen. Carrying this out, we Ðnd power-law

9 934 LAMING ET AL. Vol. 518 indices of 2È3 for R and 3È4 for R and R, with most abundance ratios being 2 insensitive to 3 the precise 4 value. Applying these corrections to the abundance ratios, in Table 3 we derive e ective FIP enhancements, which are given in Table 4. These are the ratio of the derived abundance ratio to the photospheric abundance ratio (as tabulated by Feldman et al. 1992). No entries in Table 4 are given for ratios involving the Ne VIII line. The hightemperature behavior of the G(T ) function for Ne VIII (and other Li-like ions) makes it very sensitive to the assumed power law for the emission measure, and so we neglect it in further quantitative considerations. The uncertainties quoted in Table 4 are the statistical uncertainty in the measurement of the intensity ratio, added in quadrature to the variation in the inferred ratio coming from varying the parameter a by ^1. In positions R and R the measurement errors dominate. Systematic 1 errors 2 in the atomic physics theory have been left out, since they do not a ect the observed spatial variation of a particular abundance ratio. These ratios are broadly consistent with a coronal abundance set to an accuracy of less than 30%, which is a typical uncertainty in the various atomic physics parameters required to infer abundance ratios from intensity ratios (see scatter in emission measure distribution in Figures 1È4 in Laming et al. 1995). The main exceptions are Na VII/Ne VII, where at a position away from the limb a larger FIP enhancement is seen, and the Mg VI/Ne VI ratio measured at R. The electron density derived at R by the Mg VI ratio is 2 2 also rather large, so we speculate that an undetected blend with the Mg VI line is present, or that the known S III blend is inaccurately accounted for. Something of a diminishing of the magnitude of the FIP e ect is seen in the coronal lines (discussed in more detail below). For the time being, we note that although the uncertainties in these line ratios close to the solar limb are rather large, we can be conðdent in this statement since the FIP e ect in these lines at such locations is well studied in the EUV wave band (see Laming et al. 1995), where problems of spectrum contamination by colder emission are much less severe. It is an important feature of our study that we have selected as unremarkable a region of the solar corona as possible, in order to facilitate comparison with earlier work. 4. DISCUSSION The upper transition region line ratios (Mg VI/Ne VI, Mg VII/Ne VII,NaVII/Ne VII, and Al VII/Ne VII) clearly show coronal abundances o -limb, consistent with the previously discussed notion that the UFSs are low-lying structures with quasi-photospheric abundances. The more quantitative question of how low-lying ÏÏ cannot be answered, since the two of these ratios measured on the disk (Mg VII/Ne VII and Al VII/Ne VII) give conñicting results. Close to the solar limb, we do not detect any plasma exhibiting photospheric abundances, except in the Mg VII/Ne VII ratio measured at or near the limb (see Table 4). For the most part, this is because of our inability to accurately identify and measure the upper transition region lines against the chromospheric continuum seen on the limb. In this respect an observation to image this change between abundance patterns is probably best done with the CDS instrument on SOHO, since the spectrum at the shorter wavelengths will be less contaminated by colder lines and continuum on the solar limb. Some of the other ratios listed (e.g., Mg VI /Ne VI at position R ) are also possibly a ected by such measurement problems, 2 in this case probably stemming from blending on the Mg VI line. This can also be seen in the density inferred from the Mg VI / ratio in Table 2 and position R. However, there are still a number of conclusions we can draw 2 from this study Relation to Activity and Structures This study and that of Feldman et al. (1998) are the Ðrst to measure the FIP e ect in the quiet solar corona during solar minimum. It is therefore clear that the existence of the FIP e ect is not directly related to solar activity, despite its connection to coronal heating and mass supply. This conclusion was also reached in studies of element abundances in stellar coronae of varying activity (Laming et al. 1996; Drake, Laming, & Widing 1997). Feldman et al. (1998) made abundance measurements from a polar coronal hole and an equatorial streamer, both observed at higher altitudes than the data analyzed here. These authors made the assumption that the emitting plasma was isothermal and deduced relative element abundances on this basis. This approximation can be justiðed in view of the long radiative cooling time at the densities of order 108 cm~3 in these structures and the lack of obvious time variability. However, certain temperature-dependent line ratios (see Kink et al. 1997) are not in agreement with the temperatures deduced, instead indicating values a factor of 2 or more higher, though it is not yet clear how reliable these particular diagnostic line ratios are Ionization Temperature and G(T ) e In this work we were more focussed on emission lines typically associated with T \ 8 ] 105 K plasma coming from lower altitudes, so we e adopted the more common TABLE 4 OFF-LIMB FIP FRACTIONATIONS Line ratio (A ) R 1 R 2 R 3 R 4 Mg VI /Ne VI ^ ^ ^ 0.8 Mg VII /Ne VII ^ ^ ^ ^ 0.7 Na VII /Ne VII ^ ^ ^ 1.0 Al VII /Ne VII ^ ^ ^ ^ 0.6 Mg IX /S IX ^ ^ ^ ^ 0.1 Mg IX /S X ^ ^ ^ ^ 0.9 Si IX /S IX ^ ^ ^ ^ 0.1 Si IX /S X ^ ^ ^ ^ 0.8 Si XI /(S X ] Ar XII ) ^ ^ ^ 0.3

10 No. 2, 1999 OFF-LIMB BEHAVIOR OF THE FIP EFFECT 935 approach of assuming a multithermal plasma with a particular emission measure distribution. However, it is worth commenting that the emission measures deduced from colder lines (T \ 8 ] 105 K) from various ionization stages of the same e element indicate steeper and steeper rising power laws as one moves higher up in the corona, while coronal lines indicate steeper decreasing power-law emission measure distributions. Thus the emission measure becomes closer in nature to the isothermal approximation discussed above (i.e., something like a delta function), the higher one moves above the solar limb. The di erent lines displayed in Figure 2 show di erent exponential decreases in intensity with height above the limb, hence indicating di erent temperatures of formation, again justifying our approach. Recalling the good correspondence between the G(T ) functions for S VII, and Mg VII, AlVII, and Si VII, it is interesting to note that FIP abundance ratios formed between S and the other elements indicate up to an order of magnitude too much S, after converting the observed intensity ratio to an abundance ratio using equations (1)È(4). A real abundance anomaly of this sort is ruled out by other FIP ratios including S (e.g., Mg IX/S IX). The most likely explanation is that we are observing a plasma with strong time-dependent behavior. The G(T ) curves for Mg VII, AlVII, and Si VII, in common with most other ions used in this study, essentially depend only on the ionization and recombination rates for these ions, and so time-dependent ionization e ects are similar for each ion and tend to cancel out when forming line intensity ratios (see discussion in Laming 1998). The G(T ) curve for S VII, however, is strongly a ected by the temperature dependence of the excitation rate. The A transition is 2p53s3P [ 2p53p3S, and so a 2p electron 2 1 from the ground state 2p61S must be excited to the 3p level, 0 giving it a correspondingly large excitation potential compared to the other lines, all of which involve excitations and radiative decays among various levels with principal quantum number n \ 2. S6` is a neon-like ion, the abundance of which is maximized at signiðcantly lower temperatures than Mg6`, Al6` and Si6`, and so can be expected to exhibit di erent behavior in time-dependent situations. SpeciÐcally, the unusually high excitation potential of the A transition will mean that dramatically increased emission in this line would result if the plasma were in an ionizing state, with correspondingly higher temperatures, than in a steady state Relation to the Fast Solar W ind? Another interesting feature of our FIP fractionation factors is the apparent reduction in FIP enhancement at temperatures around 106 K at the highest positions observed above the limb. It is well known that the fast solar wind elemental composition resembles that of the solar photosphere more closely than that of the corona. Studies with Ulysses data Ðnd a typical FIP enhancement of 1.5 in the fast wind (Geiss et al. 1995; von Steiger 1996). The fast solar wind is presumed to have its origin in coronal holes. Thus moving higher and higher above the solar limb close to a coronal hole, one might expect to see emission from low-fip enhanced plasma in closed magnetic Ðeld region diminishing relative to quasi-photospheric plasma emission from open Ðeld structures, especially at temperatures close to the fast wind freeze-in temperature of D1.2 ] 106 K (von Steiger 1996). Taking a fast wind particle Ñux at earth of D2 ] 108 cm~2s~1 (Schwenn 1983), we can estimate the emission measure of fast wind plasma that should be visible in our observations and compare it to what we actually see. Extrapolating we Ðnd D1013 protons cm~2 s~1 at the solar surface, giving a density of D108/f v cm~3 where f is the area Ðlling factor for the fast wind, a i.e., s the fraction of a solar surface from where the wind is emitted, and v is the wind speed in km s~1. This gives a column emission measure s of 1016 v s 2 Lf l f a 2 cm~5, where f is a linear Ðlling factor along the line of sight and L is the column l depth observed through the emitting plasma. Taking f \ f 2 and L \ R (a solar radius), the emission measure a becomes l _ 7 ] 1026 cm~5. v2 f 3@2 s a For a supposedly ubiquitous fast wind (see Habbal et al. 1997) at the altitudes we observe f D 1, Wang (1994), in a a paper concerning solar wind acceleration in coronal holes, gives v D 4 km s~1 at the solar surface (somewhere in the s lower corona) in an region outside of polar plumes. Hence the observable emission measure should be 4 ] 1025 cm~5. For comparison, the emission measure deduced from the intensity of the S IX A line at position R (the highest 4 altitude observed) is D8 ] 1025 cm~5. Thus our estimates indicate that the apparent reduction in the FIP enhancement with height in lines formed at about 106 K (very close to the fast-wind freeze-in temperature) could plausibly be explained by emission from fast solar wind plasma becoming visible and reducing the overall FIP bias. In Figure 1 it is apparent that the slit position for our observations is in fact rather a long way from the boundary of the north polar coronal hole, perhaps suggesting that we are observing a fast solar wind similar to that proposed by Habbal et al. (1997), essentially ubiquitous in origin over the whole solar disk. However, we must caution that our estimate of the fast solar wind emission measure probably could be an order of magnitude uncertain in each direction, depending on assumptions made. Longer observations at higher altitudes should provide a more deðnite answer Conclusions In conclusion, then, we have inferred that the usual FIP e ect is present in a region of di use quiet solar corona observed during solar minimum. Our study is consistent with the results of Feldman et al. (1998) coming from observations made higher up in the solar corona, but our attempts to track the behavior of element abundances all the way back to the solar limb was less successful because of the high degrees of time variability and contamination by colder lines and continuum encountered there. We anticipate that progress on this issue will come only from statistical studies assessing a large amount of data, where it can be hoped that the e ects of time variability will largely cancel out and larger regions of the spectrum can be analyzed to understand the colder emission more completely, and hence isolate the spectral regions of interest.

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