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The Astrophysical Journal, 578: 503-514, 2002 October 10
# 2002. The American Astronomical Society. All rights reserved. Printed in U.S.A.

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STRINGENT X-RAY CONSTRAINTS ON MASS LOSS FROM PROXIMA CENTAURI Bradford J. Wargelin and Jeremy J. Drake
Smithsonian Astrophysical Observatory, 60 Garden Street, Cambridge, MA 02138; bwargelin@cfa.harvard.edu, jdrake@cfa.harvard.edu Received 2002 April 5; accepted 2002 June 12

ABSTRACT We have analyzed data from two Chandra imaging observations of Proxima Centauri, searching for an Xray halo arising from charge exchange between highly charged ions in its stellar wind and neutral gas in the surrounding interstellar medium. Based upon our model of Proxima Cen's charge exchange emission, the absence of any detectable charge exchange signal places a statistical 3 upper limit of $3 Ò 10Ð13 M yrÐ1 _ _ _ (14 M ) on the mass-loss rate (9 M for the 2 limit and 4 M for 1 ), with a model uncertainty of roughly a factor of 3. This is orders of magnitude smaller than the upper limits that have been placed on late-type dwarf stars using radio observations, and it supports a recent mass-loss result for Proxima Cen based on Ly absorption profiles. We have also studied the coronal spectrum, both in quiescence and during a prominent flare. Results are consistent with those obtained in previous X-ray observations, but a firm determination of coronal metal abundances remains elusive. Subject headings: stars: coronae -- stars: individual (Proxima Centauri) -- stars: late-type -- stars: mass loss -- X-rays: stars On-line material: color figures and XMM-Newton spectra have indicated that in very active stars, the abundances of Ne, Ar, and possibly other highFIP elements can be strongly enhanced relative to that of Fe (Drake et al. 2001; Brinkman et al. 2001; Phillips et al. 2001; Maggio et al. 2002). Proxima Cen's stellar wind is also of great interest, since even a moderate M dwarf mass-loss rate would have important implications for the origin of cosmic-ray seed particles, as well as heavy-element dispersal, kinetic heating, and ionization throughout the interstellar medium (ISM). Mass loss is also of critical importance in models of angular momentum loss and other aspects of stellar evolution. With very few exceptions, however, existing measurements of stellar mass-loss rates do not extend below a few times 10Ð10 M yrÐ1 (4 orders of magnitude higher than the solar rate of $2 Ò 10Ð14 M yrÐ1) and only apply to high-mass O and B stars, red giants, and supergiants. Lim & White (1996) provide a good summary of theoretical and observational constraints on dwarf star mass-loss limits prior to 1997. A fairly recent theoretical effort is that of Badalyan & Livshits (1992), who predict that magnetically saturated M dwarf flare (dMe) stars can have winds with mass-loss rates exceeding $10Ð11 M yrÐ1. Perhaps the strongest observational evidence in favor of strong late-type dwarf winds comes from V471 Tau, an eclipsing binary consisting of a K2 V star and a white dwarf. Mullan et al. (1989) reported observing discrete absorption features in the UV continuum of the white dwarf and interpreted them as arising from a cool (104 K) wind with a massloss rate of 10Ð11 M yrÐ1. Lim, White, & Cully (1996a), however, argue that the detection of nonthermal radio emission from this system implies that any wind from the K dwarf must be optically thin, which is probably inconsistent with the temperature and mass-loss rate inferred by Mullan et al. More recently, Bond et al. (2001) observed transient À absorptions in the Si iii 1206 A line, which they ascribed to coronal mass ejections (CMEs). Based on the number of events and other viewing considerations, they estimate that the active K star `` emits some 100-500 CMEs dayÐ1, as 503

1. INTRODUCTION

Proxima Centauri, a flaring M5.5 dwarf, is our nearest stellar neighbor and lies at a distance of 1:3009 Ö 0:0005 pc (Benedict et al. 1999), roughly in the direction of the Galactic center (l Ì 313=94, b Ì Ð1=93). Proxima Cen's association with the binary system Cen is controversial (Matthews & Gilmore 1993; Anosova, Orlov, & Pavlova 1994), as the 2=2 separation between them corresponds to a large but poorly determined distance of at least a few thousand AU. However, the usual assumption is that Proxima Cen is a bound companion and shares the solar or greater photospheric metal abundances of Cen AB. A rotation period of 83:5 Ö 0:5 days, or perhaps half that, has been reported by Benedict et al. (1998), while Jay et al. (1996) have more tentatively suggested a 30 day period. Whichever period is correct, Proxima Cen is a slow rotator, consistent with its modest X-ray activity level, which is roughly onequarter the saturation limit of LX =Lbol $ 10Ð3 for active late-type stars (Agrawal, Rao, & Sreekantan 1986; Fleming et al. 1993). In addition to having a surprisingly mysterious nature, Proxima Cen is intriguing as a representative of that most common class of fully fledged hydrogen-burning stars, the M dwarfs, and it is one of the few late M dwarfs close enough to allow characterization of its coronal emission, of which coronal metallicity is a particularly interesting parameter. The reviews of, e.g., Feldman & Laming (2000), Drake (2002), and White (1996) discuss evidence from EUVE, ASCA, and BeppoSAX spectra of active stars that indicates their coronae to be significantly metal-poor with respect to a solar composition, while Drake (2002) also presents evidence that less active stars can share the solarlike coronal `` FIP effect,'' in which the coronal abundances of elements with low first ionization potentials (FIP d 10 eV; e.g., Mg, Si, and Fe) are enhanced (relative to photospheric abundances) compared to the abundances of elements with high FIP (e10 eV; e.g., O, Ne, and Ar). More recently, abundances derived from high-resolution Chandra


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compared to $1-3 dayÐ1 for the Sun. The K dwarf 's massloss rate associated with CMEs is at least ?5 25îÒ 10Ð14 M yrÐ1, but it may well be orders of magnitude higher if most of the silicon is in ionization states other than Si iii.'' Doyle & Mathioudakis (1991) and Mullan et al. (1992) have also reported marginal detections of millimeter radiation from several dMe stars, of which YZ CMi is the most significant example. Based on a 2.2 measurement at 1.1 mm and data in other wave bands from IRAS and the VLA,1 Mullan et al. (1992) inferred that YZ CMi has an optically thick wind and loses mass at the rate of $5 Ò 10Ð10 ` M yrÐ1. Houdebine, Foing, & Rodono (1990) recorded optical spectra of another dMe star, AD Leo, and argued that based on the observation of a CME during a flare, the total flare-related mass loss was between 2:7 Ò 10Ð13 and 4:4 Ò 10Ð10 M yrÐ1. Lim & White (1996), however, argued that existing radio data on both stars indicate the existence of nonthermal coronal emission, which must arise near the stellar surface and then propagate through an optically thin wind to be detectable. Citing the detection of nonthermal emission from YZ CMi at 327 MHz (Kundu & Shevgaonkar 1988), they derive a more model-dependent upper limit of 5 Ò 10Ð14 M yrÐ1 for a 300 km sÐ1 104 K wind, or up to $10Ð12 M yrÐ1 for a wind with temperature of $106 Kand velocity between 300 and 600 km sÐ1. Van den Oord & Doyle (1997) similarly derived a limit of no more than $10Ð12 M yrÐ1, based on existing IR and radio observations of several dMe stars. For the specific case of Proxima Cen, based on its nondetection at 3.5 cm, Lim, White, & Slee (1996b) directly placed a 3 upper limit of $7 Ò 10Ð12 M yrÐ1 on its mass-loss rate, assuming a wind with T $ 104 K and velocity $300 km sÐ1 (or $10Ð11 M yrÐ1 for T $ 106 K). In the past few years, several papers utilizing a different wind detection method have been published, based on the idea that interaction of an ionized stellar wind with neutral gas in the ISM, via proton-hydrogen charge exchange (CX), creates a `` hydrogen wall '' of warm neutral gas (Wood, Alexander, & Linksy 1996; Gayley et al. 1997; Izmodenov, Lallement, & Malama 1999). The signature of this gas is a slight excess of Ly absorption, which has been detected, at least tentatively, in seven nearby stars (Wood et al. 2001 and references therein). In four of those systems, three of which are composed of late-type, main-sequence stars, the absorption is strong enough to have permitted mass-loss rate esti_ _ mates: very roughly, M Ì 10 M for the active RS CVn _ type binary And (G8 IV-III+unknown) and 1 M for _ _ Ind (K5 V) (Muller, Zank, & Wood 2001); and M Ì 2 M ? _ for Cen AB (G2 V+K0 V) and an upper limit of 0.2 M for Proxima Cen (Wood et al. 2001). Some of the assumptions used to derive those rates, however, particularly regarding the intrinsic Ly profiles of the stars in question, are controversial, and the resulting uncertainties are not well understood. The D/H abundance ratios used in that work may also be less secure than was assumed (Vidal-Madjar & Ferlet 2002). In an earlier paper (Wargelin & Drake 2001), we described a more direct method of investigating the winds of

late-type dwarf stars, via the CX X-ray emission that results as highly charged ions in the wind, particularly oxygen, interact with neutral gas in the ambient ISM. The emission mechanism is essentially identical to that for comets, first explained by Cravens (1997), except that the neutral gas is primarily atomic hydrogen rather than water vapor. To briefly summarize, when a highly charged ion collides with neutral gas, an electron can be transferred from the neutral into an excited energy level of the wind ion, which then decays and emits an X ray. Two-electron CX occurs roughly 10% of the time in wind-comet interactions (Greenwood et al. 2000), but this process is unimportant in the H-dominated (single-electron) neutral gas considered here. Soon after the first X-ray observations of comets, Cox (1998) pointed out that CX must occur throughout the heliosphere as the solar wind interacts with neutral gas in the ISM. Subsequent quantitative analyses (Cravens 2000; Cravens, Robertson, & Snowden 2001) determined that this mechanism accounts for a significant fraction of the observed soft X-ray background, in agreement with indications from ROSAT observations (Snowden et al. 1995) that roughly half of the 1 keV background comes from a `` local 4 hot plasma.'' As discussed by Wargelin & Drake (2001), this same process must occur for any star with a highly ionized wind residing inside a partially neutral region of the ISM. By searching for the resulting distinctively profiled CX emission, stellar winds with mass-loss rates not much greater than the Sun's can now be detected around nearby stars with high-resolution, large-area X-ray observatories. In this paper, we present results from Chandra observations of Proxima Cen, beginning with an overview of the data and a discussion of light curves in x 2, followed by spectral analysis of emission in quiescence and during a large flare in x 3. In x 4, we derive a sensitive upper limit to the stellar mass-loss rate, based on the observed level of CX X-ray emission, marking the first application of this wind detection method, and we conclude with a discussion of model uncertainties and prospects for reducing them.
2. THE OBSERVATIONS

1 The VLA (Very Large Array) is a facility of the National Radio Astronomy Observatory (NRAO). The NRAO is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.

Proxima Cen was observed twice by Chandra as a Guest Observer target (PI J. Linsky) between 2000 May 7 and 9. Both observations (observation IDs 49899 and 641) utilized the ACIS-S CCD array in imaging mode, for 29,856 and 19,036 s, respectively. The intention was to use ACIS-S with the High Energy Transmission Grating, but a hardware fault prevented insertion of the grating. Standard Chandra X-Ray Center pipeline products were reprocessed to level 2 using the Chandra Interactive Analysis of Observations (CIAO) software version 2.1.3, taking advantage of recent gain map improvements for the central (S3) CCD released in CALDB version 2.8. Various energy filters were then applied to the data to increase the contrast of secondary X-ray sources against the background. About two dozen extraneous sources were removed from the S3 chip. Unfortunately, because of the high counting rate of the source and the chosen CCD frame time (3.2 s), more than 95% of the events in the main peak were rendered useless because of pile-up and the associated phenomenon of grade migration, which lead to distorted spectra and nonlinear counting rates. To keep these effects at a negligible level, we had to exclude data from the central core of the source.


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Fig. 1.--Light curve for observation ID 49899, using 100 s time bins (50 s for flare detail). Zero time corresponds to 2000 May 7, 22:56:24 UT. The spatial filter is an annulus of radii 4 and 30 pixels, plus the 2 pixel-wide readout streak (radii 4 and 80 pixels, with 3 pixel streak in flare detail). Background is 0.16 counts per 100 s time bin (0.40 per 50 s in flare detail). The Q1 and Q2 time ranges were included in the composite quiescent spectral analysis, and F1 was used for the flare analysis.

Because so few events were left, we included counts from the `` readout streak '' for use in spectral and light curve analysis. The streak is an artifact of the CCD readout process, with a net exposure time per frame equal to the number of CCD rows (1024) multiplied by the time required to shift the image by one row during readout (40 ls), or 0.041 s. The readout streak exposure efficiency when using a 3.2 s frame time is therefore 0:041=?3:2 ? 0:041î Ì 1:26%, and pile-up is completely negligible. Light curves for each observation ID (Figs. 1 and 2) were extracted from the source event files at energies up to 2 keV using annuli of radii 4 and 30 pixels (200 and 1500 ), centered on the source, plus a 2 pixel-wide box to include the readout streak. Background rates were derived from the entire S3 chip, after we excluded extraneous sources, a 200 pixel radius circle around Proxima Cen, and a narrow strip around the edges of the chip. The background was statistically uniform across the chip and showed no significant temporal variability during either observation.
3. SPECTRAL ANALYSIS

Readout streaks, typically containing one or two hundred counts, were used to estimate the true inner-core counting rate in the absence of pile-up, after we excluded a 50 pixel radius circle around the main peak and accounted for the minor effects of background and X-ray events in the PSF wings. Events from the background-subtracted outer core and wings were then added to give the total corrected rates. The preflare and postflare quiescent intervals of observation

3.1. Data Extraction Based upon the light curves, we divided the observations into flare and quiescent time ranges for spatial and spectral analysis. The time ranges are marked as intervals F1, F2, and Q1-4 and are summarized in Table 1. The latter part of the flare (the 200 s between F1 and Q2) was not included in F1 in order to avoid `` contamination '' of the flare's pointspread function (PSF; see x 4.3). The 1000 s between Q3 and F2 were also not analyzed.

Fig. 2.--Light curve for observation ID 641, using the same spatial filter as for observation ID 49899. Zero time corresponds to 2000 May 9, 00:05:54 UT. Background is 0.16 counts per 100 s time bin. Q3 and Q4 were included in the quiescent spectral analysis.


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TABLE 1 Spectral Data Extraction Parameters Total Spectrum Counts 983 491 475 566 914 Estimated Background Counts 47 5 18 11 23 Corrected X-Ray Counts 19,330 10,360 7,690 13,420 16,800 Ö Ö Ö Ö Ö 1560 950 950 1130 1300 Exposure Live Time (s) 28,688 592 7,450 2,172 8,186

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Time Range Q1+Q2 ......... F1 ................. Q3 ................. F2 ................. Q4 .................

Annular Radii (pixels) 3, 4, 3, 4, 3, 30 80 35 60 40

Streak Width (pixels) 2 3 3 3 3

Corrected Rate (counts sÐ1) 0.67 17.5 1.03 6.18 2.05 Ö Ö Ö Ö Ö 0.03 2.6 0.13 0.52 0.16

Note.--Spectra from time ranges Q1-Q4 were summed into a single quiescent spectrum for fitting. The spectrum of the large flare, F1, was also analyzed. `` Corrected X-ray Counts '' is the estimated total number of X-ray events (not including the readout streak) at energies below 2 keV if there had been no pile-up or grade migration effects. The live-time fraction is 0.987 for all time intervals.

ID 49899 (Q1 and Q2) have the same count rate, whereas the analogous intervals of observation ID 641 have higher rates that differ from each other by a factor of 2. To maximize the number of counts available for spectral analysis, we adjusted the sizes of the annuli and readout streak boxes for each time interval, using as small an inner radius as possible on the annuli while keeping pile-up effects negligible. Inner radii were chosen by comparing the PSFs of low-rate and high-rate data with each other and with calibration models and by studying how spectral shapes and hardness ratios varied depending upon how much of the core was included in the extraction region. As shown in Table 1, we used an inner radius of 3 pixels for quiescent times and 4 pixels for the flares. Outer radii and readout streak box widths were chosen to include as many X-ray events as possible while keeping the relative background contribution low and were different for each time interval. Extracted spectra for each of the time ranges are shown in Figure 3. Data from Q1 and Q2 were combined, since their rates and spectra were the same. Spectral analysis was performed with the CIAO SHERPA `` fitting engine '' (Freeman, Doe, & Siemiginowska 2001), using the plasma emission code MEKAL (Kaas-

tra 1992; Liedahl, Osterheld, & Goldstein 1995). The main aim of this analysis was to constrain the coronal element abundances and to investigate the plasma parameters that characterize the larger flare event. Because of the limited number of events, data were combined from all times deemed to be free of significant flaring (Q1+Q2+Q3+Q4), in order to create a composite quiescent spectrum. For the flare analysis, we used events from interval F1. Energies in the range 0.3-2.0 keV were considered; below this energy range the ACIS-S response becomes uncertain, while at higher energies the data consist largely of background events. Parameter estimation was performed using the modified 2 statistic (Gehrels 1986) and was verified using the C statistic (Cash 1979). Results are summarized in Table 2. 3.2. Quiescent Spectrum As noted in x 1, the photospheric metallicity of Cen AB (and by plausible extension, Proxima Cen) is known to be at least as high as the Sun's. We therefore first attempted to match the composite quiescent spectrum with isothermal models corresponding to the solar photospheric composi-

TABLE 2 Summary of Model Parameter Estimation Quiescent kT (keV) 0.31 Ö 0.50 Ö 0.37 Ö 0.37 Ö 0.37 Ö 0.37 Ö 1.37 0.68 0.34 0.17 0.085 0.01 0.03 0.03 0.03 0.03 0.03 Flare kT (keV) ... ... 0:81? Ð 0:81? Ð 0:81? Ð ... 1.37 0.68 0.34 ... ...

Model Isothermal, solar abundance ............ Isothermal, abundance-free ............. Isothermal, grouped abundance .......

Metallicitya [M/H] 0 ?0 1 [M/H] = Ð1:0Ð0::05 [Fe/H] = Ð0.9 Ö 0.1 [Mg/H] = Ð0.2 Ö 0.2 [Ne/H] = Ð0.2 Ö 0.1 [O/H] = Ð0.7 Ö 0.1 [M/H] 0

Normalizationb 4.1 Ö 0.1 2.0 Ö 0.1 22 Ö 2 22 Ö 2 22 Ö 2 22 Ö 2 4.6 Ö 0.3 1.4 Ö 0.3 1.4 Ö 0.2 1.5 Ö 0.4 <3.9

Metallicitya .. . .. . [Fe/H] = Ð0:6?0:: Ð0 [Mg/H] = 0:1?0::4 Ð0 3 [Ne/H] < 0.8 .. . [M/H] 0

Normalizationc ... ... 11?7 Ð4 11?7 Ð4 11?7 Ð4 ... 92 Ö 14 16 Ö 9 8.5 Ö 7.2 ... ...

5 2

0: 0: 0: 0: 0: 0:

4 1 4 1 4 1

Multithermal, solar abundance ........

.. . .. .

a All abundances are expressed in the usual logarithmic bracket notation relative to solar photospheric abundances tabulated by Anders & Grevesse 1989. See also x 3.2. b Normalizations for fits to the quiescent spectrum must be multiplied by 19.2 to account for the loss of events due to pile-up. The adjusted norR malizations correspond to the plasma EM in units of 10Ð19 =?4D2 î ne nH dV , where D is the distance to Proxima Cen (1.30 pc) and ne and nH are the electron and hydrogen number densities, respectively. c For the flare fits, normalizations must be multiplied by 21.1 to correct for pile-up losses.


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Fig. 3.--Extracted source spectra for each time segment and the background spectrum. The relative contribution of the background (dotted curves) varies in each source spectrum because of the different extraction areas and source rates for each time segment.

tion tabulated by Anders & Grevesse (1989). It should be noted that this abundance tabulation has been superseded by subsequent compilations of Grevesse and coworkers (see, e.g., Grevesse & Sauval 1998), although the differences are generally small (d0.1 dex) for the abundant elements relevant to our study (N, O, Ne, Mg, Si, S, Ar, and Fe). Two exceptions are worthy of remark here: first, Fe, for which Grevesse & Sauval (1998) adopt recent solar photospheric results that are in agreement with the value obtained from carbonaceous chondrites, ÍFe=H Ì 7:50, instead of ÍFe=H Ì 7:67; second, O, for which there is a recent solar measurement based on a non-LTE analysis of forbidden O i lines indicating an abundance of ÍO=H Ì 8:69 Ö 0:05 (Allende Prieto, Lambert, & Asplund 2001)--0.24 dex lower than the Anders & Grevesse (1989) value. While noting these differences between currently accepted values of solar photospheric abundances and those of Anders & Grevesse (1989), we retain the latter for our model analysis partly for

convenience (this compilation is `` hard-wired '' into the SHERPA fitting engine), but also to easily enable crosscomparison of our fitting results with those of other studies. As will be seen, it turns out that these abundance differences are in any event of little consequence for our modeling. It was readily apparent that isothermal models were inadequate for representing the data in the vicinity of 1 keV, below 0.5 KeV, and near 0.65 KeV (Fig. 4a). The sense of the latter discrepancy is a model flux that is significantly higher than the data indicate; the peak in the model here is largely caused by the resonance Ly line of H-like O. Allowing the global metal abundance parameter to vary yielded better matches (Fig. 4b), although with systematic problems near 0.6 keV. The best fit, although still poor, in this case was for a metal abundance relative to the solar photosphere of ÍM=H Ì Ð1:0. The plasma temperature parameter was also significantly different from that for the solar photospheric abundance case: 0.5kT (5:8 Ò 106 K),


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Fig. 4.--Composite quiescent spectrum, best-fit model, and residuals using the MEKAL radiative loss model with various assumptions. (a) Fixed solar photospheric abundances ([M/H] = 0.0). (b) Global metallicity scaling allowed to vary. The best-fit model corresponds to a metallicity of [M/H] = Ð1.0. (c) Grouped element abundances allowed to vary independently. The most important elements in the four groups are O, Ne, Mg, and Fe. The best-fit model has Ne and Mg abundance parameters somewhat higher than that for Fe (see text). (d ) Multithermal model with fixed metallicity, [M/H] = 0.0. The optimum EMs of the different temperature components are illustrated in Fig. 5.

versus 0.3kT (3:5 Ò 106 K) for the latter. An improved fit (Fig. 4c) was obtained by varying the abundances, which we grouped together because of the limited number of counts: C, N, and O; Na, Mg, Al, Si, and S; Ne and Ar; and Fe, Ca, and Ni. The most likely temperature for this case, 0.37kT (4:3 Ò 106 K), is similar to that obtained with fixed solar abundances (0.31kT ). Note that in this case, the abundance parameters for the Ne and Mg groups appear to be somewhat higher than those for the Fe and O groups (see Table 2); this is further discussed below. While an isothermal model with varying element abundances matches our data within statistical uncertainties, we also investigated the propriety of multithermal models with restricted abundance parameters. We find that a model composed of four isothermal components on a regular logarithmic temperature grid (log T Ì 6:3, 6.6, 6.9, and 7.2), all with solar photospheric abundances, also provides an acceptable match to the observations. The emission measure (EM) parameters of the different components for the optimum match are illustrated in Figure 5, together with the EMs for the different isothermal models.

Fig. 5.--ACIS-S EM parameters for the best-fit multithermal models of flare and quiescent spectra, together with individual EMs from the various isothermal models. Error bars were estimated using the full error projection utility in the CIAO tool SHERPA.


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The flare interval F1 contains fewer events than the period of quiescence, so data were grouped to yield a minimum of 10 counts in a single bin. Because F1 includes data from the flare rise, peak, and much of the decay, we would not expect the spectrum to be well fitted with a single temperature. Nevertheless, an isothermal model with varying abundance parameters (grouped as for the quiescent spectral fits) is found to be acceptable within the fairly large uncertainties of the data, although an isothermal model with a solar photospheric abundance parameter is not. The acceptability of the first fit, however, is likely just a reflection of the poor statistical quality of the data, and we hesitate to draw any firm conclusions regarding abundances. One could argue, however, that if the source were indeed isothermal, then there is evidence of a higher Fe abundance during the flare than during quiescence. Again, however, a multithermal model (log T Ì 6:6, 6.9, and 7.2) with solar photospheric abundances was found acceptable. Adding another component at log T Ì 7:5 did not yield any significant improvement to the fit. Results for all three model fits--isothermal with solar abundances, isothermal with variable abundances, and multithermal with solar abundances--are summarized in Table 2. 3.4. Discussion It is clear from the formal spectral analysis that the decimation of photon events resulting from pile-up compromises the data to an extent that useful formal constraint of model abundance parameters is not possible for either quiescent or flare emission. However, in both cases, insistence on an isothermal source model would imply that the Fe abundance is considerably lower than the solar coronal value and that Ne in particular is likely enhanced relative to Fe. Such a pattern would be in keeping with the recent results based on Chandra and XMM-Newton high-resolution spectra of the RS CVn type binaries HR 1099 (V711 Tau) and II Peg, in which strong Ne enhancements over Fe were uncovered (Brinkman et al. 2001; Drake et al. 2001; Phillips et al. 2001; Maggio et al. 2002). Moreover, Drake et al. (2001) found from a literature survey that parameter estimation analyses of low-resolution ASCA spectra of many other active stars tended to yield abundance parameters for Ne significantly higher than those for Fe, suggesting that the phenomenon is universal in active stars. Unfortunately, the quiescent data are also consistent with a multithermal model with photospheric abundances. We have no observational evidence to favor one solution over another, although we do remark that the multithermal model EMs tend to appear flatter as a function of temperature than might be expected based on EMs derived for other stars (see, e.g., Fig. 2 in the review of Bowyer, Drake, & Vennes 2000). We suggest, then, that the ACIS-S spectrum presents some evidence that the coronal abundance anomalies uncovered in other active stars are also shared by our nearest neighbor. In the case of the analysis of the flare interval F1, we are again stymied by poor data quality resulting from pile-up. Similar conclusions apply regarding the equal applicability of models in which either abundance parameters are allowed to vary or additional temperature components are added. In this regard, we note that at some level the method of parameter estimation using test statistics applied to a

large spectral range becomes meaningless when there are obvious systematic differences between model and data in any much narrower spectral interval: a test statistic that indicates statistical concordance between model and data (e.g., reduced 2 < 1) does not tell the whole story, and the propriety of the model becomes dependent on the energy range adopted for computation of the statistic. In particular, the variable-abundance parameter and multithermal models are both equally appropriate, based on the test statistic applied to the interval shown, but in the range 0.7-1 keV the multithermal model with solar photospheric abundances presents an aesthetically better representation of the data. 3.5. Quiescent Activity Level in the Context of Earlier Observations Haisch et al. (1998) presented a comparison of their RXTE observations of Proxima Cen with earlier measurements obtained by different satellites. These disparate data were compared via isothermal EM-versus-temperature loci. These types of curves are described in detail by Drake (1999). In brief, the EM-temperature (EM-T ) locus is given by the isothermal plasma EM at temperature T that yields the observed broadband count rate. The observed count rates and different responses of the various instruments that have observed Proxima Cen each define a different locus in the EM-T plane. Since, in the case of stellar coronae, the plasma is not likely to be isothermal, the individual loci actually represent the upper limit to the plasma EM at any given T. If the true EM distribution is relatively sharply peaked and the bandpasses in question are largely dominated by lines formed at this peak temperature, then the isothermal approximation can be reasonable. We have taken the EM-T loci for the different instrument count rates reported for Proxima Cen from Haisch et al. (1998) and show these in Figure 6. One small difference between our Figure 6 and Figure 3 of Haisch et al. (1998)

Fig. 6.--Loci of isothermal EM vs. isothermal plasma temperatures that correspond to different instrument count rates reported for Proxima Cen in the literature for observations between 1979 and 1996. The diamond corresponds to the isothermal EM and temperature found in the optically thin, collision-dominated model parameter estimation analysis for the case of solar photospheric abundances. Here the units of EM have been converted 2 to the logarithm of the product Ne ?T îV ?T î in units of cmÐ3. The arrow indicates the isothermal plasma temperature obtained by Haisch, Antunes, & Schmitt 1995 from their ASCA observation.


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concerns the Einstein locus, which we represent here (shaded region) as including the range of apparently quiescent count rates 0.1-0.3 counts sÐ1 from the 1979 March 6-7 and 1980 August 20 observations (Haisch et al. 1980,