THE QUIESCENT CORONA OF VB 10 Thomas A. Fleming. Mark S. Giampapa. and David Garza

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1 The Astrophysical Journal, 594: , 2003 September 10 # The American Astronomical Society. All rights reserved. Printed in U.S.A. THE QUIESCENT CORONA OF VB 10 Thomas A. Fleming Steward Observatory, University of Arizona, Tucson, AZ 85721; taf@as.arizona.edu Mark S. Giampapa National Solar Observatory, National Optical Astronomy Observatories, 1 Tucson, AZ 85726; giampapa@noao.edu and David Garza Steward Observatory, University of Arizona, Tucson, AZ 85721; locodave@hotmail.com Received 2003 March 26; accepted 2003 May 14 ABSTRACT We present results from a Chandra ACIS observation of the M8 dwarf star VB 10, a star near the hydrogen-burning mass limit. Until now, VB 10 has only been detected to flare at X-ray wavelengths. We can now report that nonflare, quiescent, X-ray emission has been detected with a luminosity, L X ¼ ð2:4 0:05Þ10 25 ergs s 1 and logðl X =L bol Þ¼ 4:9. This is consistent with the previous ROSAT nondetections of quiescent emission from VB 10. We discuss the implications of this discovery for the nature of coronae in ultracool dwarfs. Subject headings: stars: coronae stars: low-mass, brown dwarfs X-rays: stars 1. INTRODUCTION In the study of the origin of magnetic activity in the latetype stars, the M dwarfs span a critical regime in physical parameter space. It is in this region of the H-R diagram that the partially convective stellar interior becomes fully convective, with potential consequences for the nature of the operative dynamo (Rosner 1980; Giampapa & Liebert 1986; Durney, De Young, & Roxburgh 1993; Hawley, Reid, & Gizis 2000). Furthermore, the action of magnetic fieldrelated, nonradiative heating processes is most clearly manifested as a result of the vivid contrast between the cool photospheres and the overlying, hot plasma of the corona. Adopting a static loop model approach to investigate the properties of the coronae of M dwarfs, Giampapa et al. (1996) found that the two principal thermal components of the X-ray corona were characterized as follows: 1. The soft or low-temperature component, in the temperature range of T L ð2 4Þ10 6 K, consists of compact loops with lengths compared to the stellar radius of l5 R and base pressures similar to what is inferred for quiescent solar active regions. 2. The hard or high-temperature component, with temperatures T H 10 7 K, requires loop model solutions with either small filling factors (0.1), large loops (ler), and high base pressure (p 0 ep ), or very small filling factors ( f 5 0:1), small loops of length ldr, and very high pressure (p 0 4p ). These properties are reminiscent of compact flaring structures. Thus, the coronal geometry for M dwarfs appears to be dominated by a combination of relatively compact, quiescent loop configurations and an unstable flaring component. 1 The National Solar Observatory and the National Optical Astronomy Observatories are each operated for the National Science Foundation by the Association of Universities for Research in Astronomy. 982 From an observational perspective, the hard component contributes a relatively larger fraction of the total X-ray emission from the active (dme) dwarfs while also giving rise to the variability in the X-ray light curve. By contrast, in the inactive (dm) dwarfs, the soft component dominates the X-ray emission measure with little or no variability seen in the light curve. The discovery of L and T dwarfs over the past few years has motivated the question of whether these low-mass, ultracool objects can generate magnetic fields to heat coronae as efficiently as M dwarfs. All indications to this point have been that they do not. All X-ray detections of stars later than type M7 (i.e., later than the star VB 8) have been transient in nature (Fleming et al. 2000; Rutledge et al. 2001; Schmitt & Liefke 2002). While these objects were observed to flare, there was very little indication of any stable, quiescent X-ray emission similar to what is seen in the quiet solar corona or in the coronae of hotter M dwarfs. Using ROSAT, Fleming et al. (2000) measured an upper limit on the quiescent X-ray emission of VB 10 (M8 Ve; Kirkpatrick et al. 1995), which was at least 2 orders of magnitude less than that of hotter M dwarfs. This led them to suggest that nonflare, quiescent coronal emission, analogous to what is seen in solar active regions, may not occur in ultracool dwarfs such as VB 10. The amplitude of random motions at the footpoints of magnetic loops may be insufficient to jostle or stress magnetic fields. In addition, there are few ions in the very cool photosphere to couple the magnetic field with the surrounding photosphere that, in turn, is dominated by molecules. In view of the effectively neutral conditions that exist in the cool, dense photosphere of VB 10, the electrical conductivity is so low that any current system rapidly decays, and the magnetic field is in its minimum-energy configuration. Thus, the buildup and storage of magnetic field energy is excluded, and no magnetic field-related heating occurs. A more detailed treatment of this effect is given by Mohanty et al. (2002). In this scenario, the flare event would arise

2 QUIESCENT CORONA OF VB from more complex magnetic topologies that do not necessarily have footpoints in the cool photosphere, or the flares occur as a result of reconnection events in the photosphere. This hypothesis was reinforced by the detection with Chandra of a flare event in the young brown dwarf LP combined with the nondetection at a sensitive upper limit of its nonflare, quiescent emission (Rutledge et al. 2001). We note that Berger et al. (2001), using the Very Large Array, detected a radio flare in this brown dwarf as well. We are particularly interested in establishing the coronal properties of VB 10 since it is an ultracool dwarf that is accessible to sensitive X-ray observations, thus enabling us to investigate the above theory in more detail. As we report herein, our detection of quiescent X-ray emission in VB 10 is in disagreement with this hypothesis. 2. OBSERVATION We observed VB 10 with Chandra, using the ACIS backilluminated chip S3, on 2000 July 4 at UT 20:40 until 2000 July 5 UT 1:17. An X-ray source containing 27 counts was clearly detected at the Epoch position of VB 10. An extraction box of size 2>46 2>46 was used to extract the source counts within an area of 6.05 arcsec 2. The measured background count rate for the observation was 1: counts s 1 arcsec 2. We would expect the extraction box to contain 0.93 background counts after 12.6 ks of exposure time. So it is reasonable to assume that one of the 27 counts in the detection comes from the X-ray background. Therefore, we have calculated the source count rate assuming 26 source counts, although we include all 27 counts in the light curve and spectral fitting analyses described below since we have no way of knowing which count is due to the background. We analyzed the light curve of VB 10 (Fig. 1) in order to determine if the source was constant or if the light curve Fig. 1. X-ray light curve of the star VB 10 from our ACIS-S observation. The bin size is 20 minutes. showed a significant flare event. This was accomplished by using Monte Carlo methods to determine the statistical likelihood that any particular point on the light curve was constant. This particular Monte Carlo simulation generated random photon arrival times for 27 photons within an exposure time of 12.6 ks as would be expected for a constant source. Next we binned the arrival times and calculated the rate for the first and last bin. The reason that only the first and last bins are calculated is to reduce computation time. This is valid because the arrival times are homogeneous, and each bin size is the same except for the last bin, which is computed separately. Another 10,000 light curves were generated in this fashion and their rates were recorded. Once the rates have been generated and recorded, the frequencies of all the rates that occurred were determined and displayed as a histogram. The frequency of the rates when graphed looked similar to a Gaussian curve. The peak of this curve is the most probable rate that would be expected for a constant X-ray source with the given parameters. The wings of the function show the probability of various other rates, respectively. Finally, this histogram was used to determine the probability that the actual Chandra data are constant. This is accomplished by summing the probability under the Gaussian curve that was created with the specified parameters. The lower bound of integration is the lower bound error bar of a Chandra light-curve data point, while the upper bound is the upper error bar. Therefore, a data point with large errors can have a high probability of being constant because the lower and upper error bars cover the entire area that would be covered by a constant source. Also, in the case of a point that appears to be high with a small error bar, its probability would be lower because not much of its error lies in the region of rates that would be considered constant. Another case would be if the point falls on the most probable rate and has small error bars, it would have a high probability since it includes the most probable rate, and the error bars cover other rates with relatively high frequencies. This allows us to determine if there is a major flare event in a light curve in a more quantitative sense. By using binning sizes of 15, 20, and 25 minutes, we determined that all of the data points were constant with a relatively high confidence. On average, the points were above 50% in all three binning resolutions (Table 1). We determined the X-ray flux and luminosity of the source by fitting a single component Raymond-Smith plasma model to the ACIS pulse-height distribution (Fig. 2) and integrating under the curve from 0.2 to 1.5 kev. This procedure yields an X-ray flux of 6: ergs cm 2 s 1 with a corresponding luminosity of ð2:4 0:46Þ10 25 ergs s 1 and logðl X =L bol Þ¼ 4:9, where we have used a distance to VB 10 of 5.78 pc and L bol ¼ 1: ergs s 1, respectively (Fleming et al. 2000). This may be compared to the ROSAT upper limit of L X < 1: ergs s 1 and logðl X = bol Þ < 5:0, as given by Fleming et al. (2000). We barely missed detecting the quiescent coronal emission of VB 10 with ROSAT. The X-ray parameters that resulted from our analysis are summarized in Table 2. It is evident from the ACIS pulse-height spectrum shown in Figure 2 that VB 10 is a soft source with nearly all the source counts below 1 kev. The softness of the source is consistent with the coronal characteristics of a quiescent dm (i.e., nondme) star (Giampapa et al. 1996) and represents further

3 984 FLEMING, GIAMPAPA, & GARZA Vol. 594 TABLE 1 Probabilities That the VB 10 X-Ray Source Is Constant for Various Binnings of the Light Curve Probability That Source Is Constant Bin 15 minutes 20 minutes 25 minutes evidence for the predominantly nonflare origin of the X-ray emission. 3. THE QUIET CORONA We now display in Figure 3 the revised L X =L bol versus M v diagram that includes the Chandra detection of VB 10 and the upper limit for the nonflare X-ray emission in the brown dwarf LP The quiescent X-ray emission of VB 10 is 2 orders of magnitude less than that of more massive M dwarfs. Although the X-ray data are sparse at faint magnitudes, inspection of this diagram strongly suggests the onset of a sharp decline in quiescent (nonflare) coronal X-ray emission following the spectral type of VB 8 (M7 V) and continuing into the ultracool dwarf regime. Our tentative conclusion based on the X-ray data alone is nevertheless consistent with the more extensive compilation of H observations that exhibit a rather precipitous decline in H emission strength from the late M (M8 V) dwarfs through the L and T dwarf regimes (Hawley, Gizis, & Reid 1996; Gizis et al. 2000; Burgasser et al. 2002). Our assertion with respect to the X-ray emission is based on the assumption that the correlation between X-ray luminosity and H emission seen in earlier, active M dwarfs is also present among the ultracool dwarfs. For example, Fleming (1988) found that the mean value of the ratio of X-ray luminosity to H luminosity in his sample of dme stars is L X =L H ¼ 6:7 (also Fig. 2. ACIS pulse-height spectrum for VB 10, binned into 0.25 kev bins. Virtually all of the photons detected have energies less than 1 kev. Our best-fit Raymond-Smith plasma model is also shown. see Reid, Hawley, & Mateo 1995 for similar results for M dwarfs in the Hyades and Pleiades clusters). Whether this facet of the chromospheric and coronal energy balance persists in the ultracool dwarfs is not known in view of the paucity of both X-ray and H data for a similarly large sample of stars. In the specific case of VB 10, we infer an observed H emission flux of f H ¼ð4:9 9:8Þ10 15 ergs cm 2 s 1, as based on the reported range of variation in H equivalent width combined with spectrophotometry. This may be compared to our Chandra measurement of f X ¼ 6: ergs cm 2. Thus, the levels of H and X-ray emission are comparable in this object, with f X =f H 0:6 1:2. By comparison, Tinney & Reid (1998) reported for the brown dwarf LP a value of logðl H =L bol Þ¼ 5:6. This may be compared to the nonflare X-ray upper limit of logðl X =L bol Þ < 5:7 (Rutledge et al. 2000). While coronal and chromospheric flux measurements are scant for the lowest mass stars and brown dwarfs, the detection of H emission may prove to be a useful guide to the X-ray properties of ultracool dwarfs. If this is confirmed, then the H data in the literature implies, as we state in this work, that coronal X-ray emission also undergoes a rather precipitous decline from the late M (M8 V) dwarfs through the L TABLE 2 X-Ray Parameters from Recent Observations of VB 10 Parameter Exposure Time (ks) Count Rate (counts s 1 ) f X (10 13 ergs cm 2 s 1 ) L X (10 26 ergs s 1 ) logðl X =L bol Þ Chandra ACIS-S observation: Nonflare ROSAT HRI observation: Flare (mean) Flare (peak) >0.028 >6.7 >27 > 2.8 Nonflare < <0.042 <0.17 < 5.0

4 No. 2, 2003 QUIESCENT CORONA OF VB Another possibility is that as the emerging interior fields enter cool regions where the conductivity is vanishing, possible multipole moments (smaller spatial scales) will be significantly attenuated, declining as r ðnþ1þ from the region where the currents are flowing. Thus, when the surviving field finally emerges, it is the larger spatial scales (lower multipole moments) that dominate the structure of the corona in ultracool dwarfs such as VB 10. A similar argument has been invoked to explain the predominantly dipole structure of the Earth s magnetic field (Parker 1979). Fig. 3. Ratio of X-ray to bolometric luminosities plotted vs. absolute visual magnitude for M dwarfs later than type M5. The unlabeled data points are taken from Fleming et al. (1993) for nearby stars. and T dwarf regimes. On the basis of our results for VB 10, we would further argue that the presence of detectable H emission in an ultracool dwarf indicates the existence of coronal X-ray emission at a level of f X ef H (also see Cram 1982). Upon fitting the binned pulse-height spectrum, we estimate a coronal temperature of kt ¼ 0:24 0:02 kev, or T ¼ð2:80:2Þ10 6 K. We note, parenthetically, that this temperature is similar to that of the Sun-as-a-star at solar maximum (Peres et al. 2000). We find a differential emission measure given by Z 4d 2 n e n H dv ¼ 2: cm 5 ; ð1þ with an error of 4: cm 5. Assuming a fully ionized plasma and adopting coronal electron densities in the range of n e cm 3, we infer from equation (1) that the quiet corona of VB 10 is characterized by relatively largescale magnetic structures with sizes, l, in the range of l 0:2 22 R. Structures of this scale are at variance with the compact, quiescent loop configurations that appear to best describe the coronal geometry in somewhat earlier dme stars; instead, dimensions of this order are reminiscent of the geometrically large loop structures that characterize the quiet Sun (Giampapa et al. 1996). On the Sun, the larger scale structures found in quiet regions, coronal holes, and so forth, are associated with the turbulent diffusion process that is, in turn, linked to the large-scale solar magnetic dynamo. Whether this kind of solar-like, classical (!) dynamo action is present in VB 10 is unclear given that a shear layer between the lower convection zone and a radiative interior is not present in this fully convective object. Alternatively, we may speculate that the qualitative similarity of the quiescent corona in VB 10 to that of the quiet solar corona implies that dynamo processes similar to that in VB 10 also occur in the Sun. That is, the two dynamo mechanisms may coexist in the Sun, while only one kind of dynamo (presumably, a distributed dynamo; see Rosner & Vaiana 1980) is present in VB SUMMARY The observation of H emission along with transient brightenings at H, X-ray, and radio wavelengths that are reminiscent of flares demonstrate that magnetic field-related activity occurs in ultracool dwarfs and brown dwarfs. We have now detected, with Chandra, nonflare, quiescent coronal emission in at least one ultracool dwarf, VB 10. The detections and upper limits in both the X-ray and at H are consistent with the onset of a real decline in chromospheric and coronal heating efficiencies (L X;H =L bol ) in the ultracool dwarf regime. We estimate the temperature and emission measure of the coronal plasma in VB 10 and find that it is a soft source. A preliminary analysis based on the inferred emission measure suggests that the coronal morphology in this object is dominated by large-scale magnetic structures, similar to that of the quiet corona of the Sun. We speculate that any possible small-scale structures that are characterized by higher multipole moments of the magnetic field are significantly attenuated as the emerging dynamo-generated field traverses relatively cool, nonconducting regions of the interior. Our detection of the counterpart of solar quiescent coronal emission in VB 10 is at variance with the hypothesis discussed by Fleming et al. (2000; see x 1), namely, that stars as cool as VB 10 do not have quiescent (10 6 K) coronal plasma owing to the low conductivity in the largely neutral photosphere of such stars, which, in turn, inhibits coronal heating to the point where quiescent, coronal plasma does not even exist. In light of the observational results presented herein, this explanation is either incorrect or it applies only to even lower photospheric temperatures (i.e., T eff < 2700 K) in the ultracool dwarf regime. The further examination of this hypothesis will require deep X-ray observations of ultracool dwarfs. Alternatively, H can serve as a useful proxy for the degree of coronal X-ray emission in an ultracool dwarf. In this context, it is important to note that coronal emission may be completely absent in those ultracool and brown dwarfs that do not have H emission in their visible spectra. We tentatively suggest that f H f X in those very low mass stars that exhibit H emission. Using this as a guide to the practical feasibility of detection in the X-ray (outside of flares) by either Chandra or XMM-Newton leads to estimates of exposure times exceeding 100 ks for L and T dwarfs at distances greater than 10 pc. We acknowledge interesting discussions with J. R. Jokipii and B. Durney that materially contributed to this investigation. We would also like to thank the referee, Neill Reid, for making suggestions that improved this paper.

5 986 FLEMING, GIAMPAPA, & GARZA Berger, E., et al. 2001, Nature, 410, 338 Burgasser, A. J., Liebert, J., Kirkpatrick, J. D., & Gizis, J. E. 2002, AJ, 123, 2744 Cram, L. E. 1982, ApJ, 253, 768 Durney, B. R., De Young, D. S., & Roxburgh, I. W. 1993, Sol. Phys., 145, 207 Fleming, T. A. 1988, Ph.D. thesis, Univ. Arizona Fleming, T. A., Giampapa, M. S., & Schmitt, J. H. M. M. 2000, ApJ, 533, 372 Fleming, T. A., Giampapa, M. S., Schmitt, J. H. M. M., & Bookbinder, J. A. 1993, ApJ, 410, 387 Giampapa, M. S., & Liebert, J. 1986, ApJ, 305, 784 Giampapa, M. S., Rosner, R., Kashyap, V., Fleming, T. A., Schmitt, J. H. M. M., & Bookbinder, J. A. 1996, ApJ, 463, 707 Gizis, J. E., Monet, D. G., Reid, I. N., Kirkpatrick, J. D., Liebert, J., & Williams, R. J. 2000, AJ, 120, 1085 Hawley, S. L., Gizis, J. E., & Reid, I. N. 1996, AJ, 112, , in ASP Conf. Ser. 212, From Giant Planets to Cool Stars, ed. C. A. Griffith & M. S. Marley (San Francisco: ASP), 252 REFERENCES Kirkpatrick, J. D., Henry, T. J., & Simons, D. A. 1995, AJ, 109, 797 Mohanty, S., Basri, G., Shu, F., Allard, F., & Chabrier, G. 2002, ApJ, 571, 469 Parker, E. N. 1979, in Cosmical Magnetic Fields: Their Origin and Their Activity (Oxford: Clarendon), 716 Peres, G., Orlando, S., Reale, F., Rosner, R., & Hudson, H. 2000, ApJ, 528, 537 Reid, N., Hawley, S. L., & Mateo, M. 1995, MNRAS, 272, 828 Rosner, R. 1980, in SAO Rep. 389, Cool Stars, Stellar Systems, and the Sun, ed. A. K. Dupree (Cambridge: SAO), 79 Rosner, R., & Vaiana, G. S. 1980, in X-Ray Astronomy, ed. R. Giaconni & G. Setti (Dordrecht: Reidel), 129 Rutledge, R. E., Basri, G., Martín, E. L., & Bildsten, L. 2000, ApJ, 538, L141 Schmitt, J. H. M. M., & Liefke, C. 2002, A&A, 382, L9 Tinney, C. G., & Reid, I. N. 1998, MNRAS, 301, 1031