X-Ray Properties of the Young Stellar And

X-RAY PROPERTIES OF THE YOUNG STELLAR AND SUBSTELLAR OBJECTS IN THE IC 348 CLUSTER: THE CHANDRA VIEW Thomas Preibisch Max-Planck-Institut fu¨r Radioastronomie, Auf dem Hu¨gel 69, D-53121 Bonn, Germany; [email protected] and Hans Zinnecker Astrophysikalisches Institut Potsdam, An der Sternwarte 16, D-14482 Potsdam, Germany; [email protected] Received 2001 October 2; accepted 2001 November 20

ABSTRACT We explore the X-ray properties of the young stellar and substellar objects in the open cluster IC 348 as seen in our deep Chandra X-Ray Observatory Advanced CCD Imaging Spectrometer image. First, we give identifications of all X-ray sources and determine upper limits for the X-ray luminosities of the undetected cluster members. Then we analyze the X-ray spectra of the young stellar objects, deriving plasma temperatures between 0.7 and 3 keV for the T Tauri stars in IC 348 and higher temperatures, between 3 and 7 keV, for flaring sources and two embedded young stellar objects. We find several large X-ray flares, in some of which a clear hardening of the X-ray spectra during the flare peak is seen. Next we use the exceptional optical, infrared, and X-ray data set of this cluster to study various correlations and their implications, and to dis- cuss new answers to some long-standing questions related to X-ray emission from young (sub)stellar objects. The X-ray luminosities of the young low-mass stars are strongly correlated to the stellar bolometric luminosi- 4 ties (LX 10 Lbol). Also, a good correlation between X-ray luminosity and stellar mass is found 2 (LX / M ). For the weak-line T Tauri stars we find a tight correlation between X-ray activity and chromo- 0:8 spheric activity ðLX / LHÞ, supporting the hypothesis that the chromosphere is heated by X-rays from the overlying corona. The observed X-ray properties of the brown dwarfs (and brown dwarf candidates) are very similar to those of late-type stars; we explain this behavior as the consequence of the fact that very young substellar objects are still warm enough to maintain partially ionized atmospheres, which are capable of sustain- ing electrical currents, while in the cooler neutral atmospheres of L and T dwarfs such currents are shut off (hence no X-ray emission). Finally, we explore the difference between the X-ray luminosity functions of classical and weak-line T Tauri stars. We find that the classical T Tauri stars in IC 348 seem to be on average less X-ray luminous than the weak-line T Tauri stars. However, we suggest that this apparent difference is caused by a selection effect: there is a strong detections bias against those weak-line T Tauri stars that are optically faint and hence X-ray faint; the population of classical T Tauri stars, on the other hand, is essentially com- pletely known because of its very prominent H emission. This conclusion is corroborated by another new result: when using a photometrically selected, magnitude-limited, complete sample of T Tauri stars and tak- ing the KL infrared excess as a tracer of circumstellar material, we find no evidence in IC 348 for a difference in X-ray properties of young stars with and without circumstellar matter, i.e., classical and ‘‘ naked ’’ T Tauri stars. Key words: low-mass, brown dwarfs — open clusters and associations: individual (IC 348) — stars: coronae — stars: pre–main-sequence — X-rays On-line material: machine-readable tables

1. INTRODUCTION Advanced CCD Imaging Spectrometer (ACIS) on board the Chandra X-Ray Observatory, in which we detected IC 348 is a very young stellar cluster in the Perseus molecular cloud complex. The cluster is exceptionally 215 individual X-ray sources. First results of this observation have been reported in Preibisch & Zinnecker well studied in optical and infrared wavelengths (see Her- (2001, hereafter PZ01), including a comparison of the big 1998, hereafter H98; Scholz et al. 1999; Lada & Lada optical, infrared, and X-ray images. In this paper we 1995; Luhman et al. 1998, hereafter L98; Luhman 1999, present detailed information on the identification of the hereafter L99; Najita, Tiede, & Carr 2000, hereafter X-ray sources and study the X-ray properties of the N00). Spectral types are known for more than 200 stars, young stellar and substellar objects in IC 348, including H emission has been detected for more than 100 stars, their implications. and numerous substellar objects have been identified. The distance to IC 348 is 310 pc (H98), the mean age of the stars is 1.5 Myr, and their mean extinction is 3.5 mag (H98; L98). ROSAT X-ray observations of IC 348 were 2. Chandra Observations and Data Analysis presented by Preibisch, Zinnecker, & Herbig (1996, here- The Chandra observation of IC 348 was performed on after PZH96) and led to the discovery of 116 X-ray 2000 September 25, utilizing ACIS in its imaging configura- sources in a 2 diameter field of view. We have recently tion. The total exposure time was 52,956.8 s. The results pre- obtained a deep X-ray image of IC 348 with the sented in PZ01 can be summarized as follows: 1613 1614 PREIBISCH & ZINNECKER Vol. 123

TABLE 1 Chandra X-Ray Sources in IC 348

Source Identiﬁcation SpT AV W(H) kT LX CXOPZ- (2) (3) (mag) (A˚) (keV) (1028 ergs s1) (1) (4) (5) (6) (7)

J034346.9+321320 ... 2MASS ... 0.70 i ...... 5.0 1.6 J034349.3+321040 ... H-IfA68 ... 1.74 i 2.0 ... 34.2 7.0 J034351.2+321309 ... L-22 G5 2.31 ... 2.1 0.3 472.0 40.3 J034351.5+321239 ... noCP ...... J034353.3+320927 ... noCP ......

Notes.—Table 1 is presented in its entirety in the electronic edition of the Astronomical Journal. A portion is shown here for guidance regarding its form and content. Columns: (1) Chandra X-ray source name as defined in Preibisch & Zinnecker 2001; the name contains the J2000.0 coordinates in the format HHMMSS.S+DDMMSS. (2) Identification with optical and/or infrared counterparts from the catalog of Herbig 1998 (H-), Luhman et al. 1998 or Luhman 1999 (L-), or Najita et al. 2000 (N-). Stars known to be foreground or background objects are marked with ‘‘ FG ’’ and ‘‘ BG ’’, respectively. If no cataloged counterpart is known, the entry ‘‘ 2MASS ’’ means that the X-ray source has a counterpart in the 2MASS point-source catalog, while ‘‘ noCP ’’ means that no optical or infrared counterpart could be found. Source CXOPZ-J034424.6+321349 is identified with a faint infrared source not included in the 2MASS catalog. (3) Spectral type of the counterpart. (4) Visual extinction of the counterpart; ‘‘ i ’’ indicates that the value was estimated from the near-infrared colors, and a colon indicates that, owing to a lack of further information, an extinction of 3.5 mag has been assumed. (5) Equivalent width of the H emission line. Negative values denote H absorption. (6) Plasma temperature determined in the fit of the X-ray spectrum. (7) Extinction corrected 0.2–10 keV band X-ray luminosity.

1. The number of individual X-ray sources detected was be rather uniform, we are confident that they are back- 215. ground objects. 2. About 80% of all known cluster members with masses In Table 1 we present a list of all X-ray sources with corre- between 0.15 and 2 M are seen as X-ray sources in our sponding identifications and further information about the Chandra image. stellar and X-ray properties. The determination of the 3. X-ray emission at levels of 1028 ergs s1 was discov- extinction corrected X-ray luminosities (integrated over the ered from four of 13 known brown dwarfs and from three of 0.2–10 keV band) was based on the analysis of the X-ray 12 brown dwarf candidates in IC 348. spectra as described in x 5. 4. X-ray emission was also detected from two deeply embedded young stellar objects, presumably Class I protostars, south of the optical cluster center. 4. UPPER LIMITS FOR NONDETECTED The data analysis presented in this paper was performed CLUSTER MEMBERS with the CIAO 2.1.2 software package provided by the In Table 2 we present a list of all those known cluster Chandra X-Ray Center and is based on the Level 2 proc- members that were not detected as X-ray sources in our essed event list provided by the pipeline processing at the Chandra image. We determined upper limits to the count Chandra X-Ray Center. rates of these objects by counting the observed number of photons in source regions centered at their optical/ infrared positions and comparing them with the expected 3. IDENTIFICATIONS OF THE CHANDRA number of background photons determined from several X-RAY SOURCES large source-free background regions. We used the Baye- We searched for counterparts of the X-ray sources in the sian statistics method described by Kraft, Burrows, & member lists published in H98, L98, L99, and N00. For X- Nousek (1991) to determine the 90% confidence upper ray sources without a cataloged counterpart, we searched limits for their count rates. From these count-rate upper for optical counterparts on the red DSS plate and on our limits we computed upper limits for the extinction-cor- own deep optical images, and for infrared counterparts in rected X-ray luminosities in the 0.2–10 keV band assum- the 2MASS source catalog, in the 2MASS images, and in ing thermal plasma spectra with a temperature of kT =1 our own deep J- and K-band images1 of IC 348. Most of the keV (see Feigelson & Montmerle 1999; Preibisch 1997; X-ray sources can be well identified with known cluster see also x 5) and computing the absorbing hydrogen col- members. Some 40 X-ray sources do not have an optical or umn density from the visual extinction according to the 21 2 infrared counterpart and are most likely background (prob- relation NH = AV 1.8 10 cm (Paresce 1984; see ably extragalactic) objects. This number is consistent with also Predehl & Schmitt 1995). the expected number of extragalactic background X-ray sources based on the observed log N log S statistics from 5. X-RAY SPECTRA the deep X-ray source counts in the Chandra Deep Field South (see PZ01 for details). Also, as the spatial distribution For all sources containing at least 200 source counts of these unidentified X-ray sources in our image appears to we performed a spectral analysis. The aim was to derive the plasma temperature for each source and to determine X-ray luminosities. The extraction of the X-ray spectra 1 The limiting magnitudes of our images are J 19 and K 19. was performed using the CIAO thread PSEXTRACT. No. 3, 2002 X-RAY EMISSION IN IC 348 1615

TABLE 2 Properties and X-Ray Upper Limits for Undetected Cluster Members

Star Other Names SpT AV W(H) Count Rate LX IC 348-J (2) (3) (mag) (A˚) (counts ks1) (1028 ergs s1) (1) (4) (5) (6) (7)

034346.4+321106... H-IfA2 ... 2.08 i 85.0 <0.11 <1.33 034348.6+321350... H-IfA3 ... 0.74 i 13.0 <0.23 <1.73 034351.5+321123... H-1 ... 3.50: ... <0.08 <1.64 034353.1+320844... H-3 ... 4.03 i ... <0.07 <1.64 034353.5+320726... H-5 ... 21.64 i ... <0.07 <17.24

Notes.—Table 2 is presented in its entirety in the electronic edition of the Astronomical Journal. A portion is shown here for guidance regarding its form and content. Columns: (1) Star name based on the J2000.0 coordinates in the format IC348-JHHMMSS.S+DDMMSS. (2) Name in the catalog of Herbig 1998 (H-), Luhman et al. 1998 or Luhman 1999 (L-), or Najita et al. 2000 (N-). (3) Spectral type. (4) Visual extinction; ‘‘ i ’’ indicates that the value was estimated from the near-infrared colors, and a colon indicates that, owing to a lack of further information, an extinction of 3.5 mag has been assumed. (5) Equivalent width of the H emission line. (6) 90% conﬁdence upper limit to the ACIS count rate. (7) 90% conﬁdence upper limit to the extinction corrected 0.2–10 keV band X-ray luminosity.

This script extracts the pulse-height spectra for a source The fitting results were also used to compute the and a background region, groups the source spectrum, extinction-corrected 0.2–10 keV band X-ray luminosities and builds the proper redistribution matrix files (RMFs) of the sources with XSPEC. In Figure 2 we show the and ancillary response files (ARFs). We then performed plasma temperatures plotted versus the X-ray luminosi- spectral fits with the Sherpa package contained in CIAO. ties. For most objects we find plasma temperatures We used the XSPEC model RAYMOND describing the between 0.7 and 3 keV, which is just the typical range for emission from a thermal plasma spectrum and used X-ray–active T Tauri stars (see Feigelson & Montmerle WABS for the absorption model. In the fitting process 1999; Preibisch 1997). For some sources we find rather model spectra are simulated, folded through the corre- high plasma temperatures, up to 6 keV; many of these sponding response files, and compared with the observed particularly hard sources showed an X-ray flare during pulse-height spectrum; the model parameters are adjusted our observation. The effect of spectral hardening during until the best agreement between the model and the flares will be investigated in detail in x 6. observed spectrum is found. For X-ray sources with less than 200 counts, for which we Although we are fully aware that the coronae of active did not extract spectra, we determined a transformation fac- stars are generally not monothermal (e.g., Brickhouse et al. tor from count rate to X-ray luminosity, assuming a plasma 2000; Sanz-Forcada, Brickhouse, & Dupree 2001), we tried temperature of 1 keV and computing the absorbing hydro- to fit all spectra with a single-temperature plasma model gen column density from the visual extinction. The assumed plus absorption. This approach is justified because the pur- temperature of 1 keV is consistent with the mean tempera- pose of our analysis was not to perform a detailed investiga- ture determined in the fits of the weaker sources. tion of the temperature structure of those sources (which is hardly possible given the rather low numbers of detected X- ray photons in each source), but because we want to charac- terize the coronal temperature distribution by a single 6. TEMPORAL VARIABILITY parameter, i.e., some kind of a ‘‘ characteristic ’’ tempera- We extracted light curves for all X-ray sources with the 2 ture, which then can be related with other stellar parame- CIAO command DMEXTRACT to study their temporal ters. For some sources, the single-temperature model did variability. Relatively short, 60 s time bins were used for the not yield an acceptable fit, i.e., resulted in a v2 probability of extraction in order not to loose information. In addition to Q < 0.05 for the best-fit model spectrum and the observed the source light curves, we also extracted the light curves spectrum to be drawn from the same underlying spectrum. from seven large source-free background regions. Light- In these cases we used a two-temperature model, which curve extraction was done for the full Chandra energy band, always gave statistically acceptable fits, and then used the as well as separately for the ‘‘ soft ’’ band (0.2–2 keV) and emission-measure weighted mean of the two temperatures the ‘‘ hard ’’ band (2–10 keV). The individual background as the characteristic plasma temperature TX. The plasma light curves were combined into a mean background light temperatures are listed in Table 1. Figure 1 shows the spec- curve, which then was subtracted from each of the source tra of 12 individual sources, which are representative for light curves. our data. A statistical evaluation of the light curves indicated significant temporal variability in about half of the light curves. Closer inspection of the individual light curves showed irregular count-rate variations in most cases. For 18 2 As demonstrated, for example, by Peres et al. (2000), fits of simulated spectra based on continuous temperature distributions with simple single- sources, however, we found clear evidence for systematic temperature models usually yield temperatures near the peak of the under- variability, mostly in the form of flares. These events will be lying temperature distribution. discussed in the next paragraph. 1616 PREIBISCH & ZINNECKER Vol. 123

Fig. 1.—Some representative X-ray spectra of young stellar objects in IC 348. The solid dots with error bars show the observed spectra, the solid histogram lines show the best-ﬁt models.

6.1. Sources Showing Strong Variability light curves display a wide range of different morphologies. Several objects show typical flares with rise and decay times To increase the signal-to-noise ratio in the light curves, of a few hours, very similar to flares on other X-ray active we rebinned them by a factor of 30 (or 20 for the strongest stars (e.g., Montmerle et al. 1983; Gagne´, Caillault, & sources), resulting in a time resolution of 30 minutes (20 Stauffer 1995; Preibisch & Neuhaüser 1995). The ROSAT minutes). In Figure 3 we present the light curves for those 18 observations of IC 348 (PZH96) already showed several sources that showed significant systematic variability. These large X-ray flares in IC 348, among them a giant X-ray flare No. 3, 2002 X-RAY EMISSION IN IC 348 1617

amplitude). The light curve of CXOPZ J034444.9+321337, on the other hand, shows significant variability in the soft- band light curve only. Another interesting feature in several light curves (in CXOPZ J034508.0+320402, J034450.7+ 321904, J034440.7+321306, J034425.5+321130, and J034422.3+321201) is that the peak in the hard-band light curve occurs about 1 hr earlier than the peak in the soft- band light curve. This effect is well known from detailed observations of solar flares (e.g., Tanaka 1986; Sato 2001; Reale, Peres, & Orlando 2001) and can be understood in the light of the magnetic reconnection flare model: the nonther- mal electrons accelerated at the coronal reconnection site collide with chromospheric plasma, resulting in strong heating of the plasma at the impact site to very high temperatures and the emission of hard X-rays. The softer X-ray emission comes from the hot plasma in the postflare loops, which are filled only after the impulsive flare phase. Several objects show light curves with shapes not resem- bling those of typical flares. CXOPZ J034416.4+320955 shows a rapid increase of the count rate without a following significant decay. CXOPZ J034425.5+321230 and Fig. 2.—Plasma temperature vs. X-ray luminosity as determined in the J034424.2+321019 show a slow continuous rise of their spectral fits for the strong X-ray sources among the young stellar objects in count rates. The light curve of CXOPZ J034426.0+3204303 IC 348. No correlation is seen. Sources that showed a large X-ray flare dur- might suggest periodic variability. These light curves may ing the observation (see x 6) are marked by triangles; these flare sources be explained by rotational variability: for example, the X- obviously have higher than average plasma temperatures. The two embedded objects (see x 8) are marked by gray circles around the solid dots ray emission may be dominated by one or several relatively (note their very high plasma temperature). small active coronal regions, which are initially on the back side of the star and gradually move to the front side due to the stellar rotation (see Stelzer et al. 1999). on the classical T Tauri star LH 92 (Preibisch, Zinnecker, & Schmitt 1993). The number of flaring events (14) seen in 6.2. X-Ray Spectral Variations during Strong Flares our Chandra data is fully consistent with the expected flare Three of the flaring sources have enough counts to allow frequency based on the ROSAT results. We have deter- the extraction of X-ray spectra in different flare phases. For mined the rise time, decay time, and the amplitude of these these objects we extracted X-ray spectra during the flare flares, and we summarize these parameters in Table 3. peak phase and during the quiescent phase before the occur- Many stellar X-ray flare observations reported in the lit- rence of the flare. These spectra are shown in Figure 4. erature have been analyzed with quasi-static cooling models We first fitted the preflare spectra with single-temperature (e.g., van den Oord & Mewe 1989) to infer some basic plasma models; this yielded the quiescent temperature Tq. parameters of the flaring plasma, such as the plasma density For the fit of the flare spectra we used a two-component and the length scale of the flaring loop. In most cases such model, in which the temperature and emission measure of & analysis yielded very large loop lengths (l R*). The quasi- the first component was fixed to the values found from the static cooling models, however, are based on the assumption preflare spectrum, while the parameters for the second com- that no (substantial) energy input occurs after the peak of ponent were allowed to vary. The fit then yields the temper- the flare. Recently, more and more evidence has accumu- ature Tf of the additional X-ray flux that can be attributed lated that this assumption is not valid. The analysis of a to the flare. We find kT = 1.0 keV and kT 6.4 keV for large X-ray flare on Algol by Schmitt & Favata (1999) and q f CXOPZ J034450.7+321904, kTq = 1.1 keV and kTf 6.2 Favata & Schmitt (1999), for which an eclipse allowed to keV for CXOPZ J034440.1+321134, and kT = 0.7 keV estimate a geometrical size for the flaring plasma, clearly q and kTf 3.7 keV for CXOPZ J034424.2+321019. The showed that the quasi-static cooling models overestimate spectral hardening during the flares can be seen well in the the loop size by a large factor. We will therefore not use comparison of the preflare spectra to the flare spectra in these models to analyze the flares observed in our data. Figure 4. More realistic flare models must be based on hydrodynamic In this context it is also interesting to consider the relation models which take into account sustained heating during between the plasma temperature and the emission measure the flare decay (e.g., Reale et al. 1997; Favata, Micela, & in the flare phase. Feldman, Laming, & Doschek (1995) and Reale 2001). Application of such models requires time-re- Shibata & Yokoyama (1999) found a good correlation solved spectroscopy during the flare-decay phase; this is between these two parameters for a wide range of different hardly possible for our flaring objects, given the relatively flaring events, from solar microflares via stellar flares to very low numbers of detected counts. powerful flares on T Tauri stars and protostars. In the plot Some of our light curves show very interesting features. of flare temperature versus emission measure shown in Ima- In several cases the soft- and hard-band light curves exhibit nishi (2001a; their Fig. 7), the parameters of the three flares pronounced differences. In source CXOPZ J034436.9+ 320645, only the hard-band light curve shows a significant increase during the flare, while the soft-band light curve 3 This is LH 92, the CTTS that showed a huge X-ray flare during the remains more or less constant (at least has much lower ROSAT PSPC observation of IC 348 (see Preibisch et al. 1993). 1618 PREIBISCH & ZINNECKER Vol. 123

Fig. 3.—Light curves for the sources exhibiting strong variability during our Chandra observation. The histograms show the full energy-band background- subtracted source light curves, rebinned into 30 minute bins (20 minute for the strongest sources). The gray solid lines shows the light curves in the soft (0.2–2 keV) band, the dark solid lines the light curves in the hard (2–10 keV) band. These light curves demonstrate the wide range of different kinds of temporal variations. in IC 348 are in very good agreement with this general corre- 348. Up to now X-ray emission has been detected from only lation. This suggests that these flaring events are a scaled-up a very few brown dwarfs (Neuhaüser et al. 1999; Rutledge et version of solar flares. al. 2000; Feigelson et al. 2002; Imanishi et al. 2001b). Thus it is appropriate to describe the X-ray properties of these objects in IC 348 in detail. The very small number (<20) of 7. X-RAY PROPERTIES OF THE LOWEST-MASS photons detected from each of these very low mass objects (SUBSTELLAR) MEMBERS prevents detailed spectral fitting, as well as extraction of a meaningful light curve. What we can do, however, is to As described in PZ01, we detected weak X-ray emission investigate the energies and the arrival times of the individ- from seven brown dwarfs and brown dwarf candidates in IC ual photons detected from each object (see Fig. 5). No. 3, 2002 X-RAY EMISSION IN IC 348 1619

TABLE 3 Strongly Variable X-ray Sources in the Chandra Data of IC 348a

Source Amplitude Rise Time Decay Time CXOPZ- SpT Type (crmax/cr0) (hr) (hr)

J034355.4+320932 ... K0 Flare 5 2 1.0 J034359.6+321403 ...... Flare 35 2 3.0 J034416.4+320955 ... K3 Step 2.8 2 >10 J034418.0+321053 ... K7 Decay 10 ... 1.7 J034422.3+321201 ... M2 Flare 9.0 1 2.0 J034424.2+321019 ... K5 Rise 8.8 2.7 ... J034425.5+321130 ... M0 Flare 14 ... 2.0 J034425.5+321230 ...... Rise 10 4 ... J034426.0+320430 ... G8 Periodic ? 3.5 6 10 J034426.6+320358 ... M4.75 Flare 7.7 2.5 6.5 J034436.9+320645 ... G3 Flare 1.8 1 0.5 J034437.8+320804 ... K7 Decay ...... 8.0 J034437.9+320329 ... K6 Flare 7 1 1.0 J034440.1+321134 ... K2 Flare 8.4 1 1.5 J034440.7+321306 ...... Flare 11 1 0.5 J034444.9+321337 ...... Flare 8.8 ... 2.2 ... Flare 6.4 ... 2.5 J034450.7+321904 ... A4 Flare 6.0 2 2.5 J034501.8+321428 ...... Flare 4.7 0.5 1.5 ... Flare 6.7 0.5 1.5

a We list the source name, spectral type, the character of the variation, the amplitude of the count rate variation, the rise time, and the e-folding decay time.

First we consider possible temporal variability. It is mine the maximum number Nmax of detected photons in any immediately apparent that for most objects the photon such time-period. The Poisson probability to find Nmax or arrival times seem to be rather uniformly distributed. Only more counts in such a period for which the expected number the brown dwarf L-105-01 is obviously variable: seven of of counts is hNi = 2 is given by the nine photons detected from this source arrived in a 3.5 NXmax1 hr period, while only two photons arrived in the remaining 9 2k P ¼ 1 e2 : hr of the observation. To get a more quantitative assessment k! of the variability, we performed a statistical evaluation of k¼0 the photon arrival times in the following way: Given the P gives the probability that the observed maximum number total number photons detected in each source region, we of counts Nmax is caused by a statistical fluctuation, while computed the mean count rate, and from this we specify the the ‘‘ probability of variability ’’ (pov) = 1P can be identi- time period in which one would expect to detect just two fied as the confidence level for the rejection of the null photons, under the assumption of a constant count rate. hypothesis that the source count rate is constant. The pov Then we used a ‘‘ sliding time-window ’’ approach to deter- values determined in this way are given in Figure 5. For

Fig. 4.—Comparison of X-ray spectra extracted during the quiescent period and during the ﬂare phase for three young stars in IC 348 1620 PREIBISCH & ZINNECKER Vol. 123

Fig. 5.—Photon arrival times and energies for the X-ray–detected brown dwarfs and brown dwarf candidates in IC 348. The quantity pov gives the Poisson probability for temporal variability. Only the object N-105-01 is clearly variable, while N-045-02 and L-312 may be variable. The photon energies are seen to be of the order of 1 keV. most objects we find no evidence for temporal variability are essentially the same as those of (older) low-mass (late- (pov < 95%). The objects N-045-02 and L-312 show weak M–type) stars. This begs the question of the degree to which evidence for variability (pov = 95%). The brown dwarf N- these findings are surprising, i.e., whether or not one should 105-01, however, is clearly variable (pov = 99.9%); the pho- expect the X-ray properties of these substellar objects to be ton arrival times are consistent with a flare occurring about different from those of low-mass stars. 5.5 hr after the start of our observation. Recent investigations of very low mass stars and Next we consider the energies of the individual X-ray brown dwarfs in the solar neighborhood have found a photons from the very low mass objects. With the exception strong change in the activity of these objects at the bot- of object L-613 the energies of the photons detected from tom of the main sequence; there is clear evidence for a these objects lie generally between 0.5 and 2 keV. Simu- drop in the chromospheric and coronal activity level at a lations of pulse-height spectra with XSPEC show that the spectral type around M9 (Gizis et al. 2000; Fleming, observed distributions of photon energies are fully consis- Giampapa, & Schmitt 2000; Martin & Bony 2001), i.e., tent with a plasma temperature of kT 1–2 keV. This is at the stellar/substellar boundary. The reason for this similar to the temperature found for most objects in the drop is probably related to the fact that these very cool spectral fits (see x 5). objects have neutral atmospheres with very high electrical Other recent results on the X-ray emission from young resistivity,4 in which the rapid decay of currents prevents brown dwarfs (e.g., Imanishi et al. 2001b; Feigelson et al. the buildup of magnetic free energy and therefore cannot 2002) have yielded X-ray properties very similar to those we provide support for magnetically heated chromospheres find here for the brown dwarfs in IC 348. These properties and coronae (see Fleming et al. 2000). can be summarized as follows: These arguments, however, are not relevant for the brown dwarfs in IC 348 (and those in other star-forming regions), 1. The fractional X-ray luminosities are in the range L / X because they are still very young. These very young brown L 104–103. bol dwarfs are of mid- to late-M spectral types and are warm 2. The typical plasma temperatures are T 1–2 keV. enough to maintain a partially ionized atmosphere. As the 3. Mostly quiescent X-ray emission is detected from photospheric properties of these young substellar objects are these objects, but some brown dwarfs show evidence for basically the same as those of older stellar objects, it is not temporal variability, including possible flares. surprising that their coronal properties are similar to those These observed X-ray properties of young brown dwarfs of low-mass stars. In other words, the very young brown No. 3, 2002 X-RAY EMISSION IN IC 348 1621 dwarfs in IC 348 ‘‘ do not yet know ’’ that they are brown dwarfs and not stars, and therefore ‘‘ behave ’’ like low-mass stars. The consideration of all available observational results suggests the following picture for the evolution of activity in very low mass objects: At ages of a few megayears the brown dwarfs still have relatively warm atmospheres. Although they are fully convective, and therefore the standard solar- like dynamo cannot work, they are somehow able to sustain a hot corona, just like the fully convective late-type stars, which are known to be strong coronal emitters. Poten- tial dynamo mechanisms are small-scale dynamo action in a highly turbulent convection zone (see Giampapa et al. 1996 and references therein) or an 2 dynamo, as suggested, for example, by Ku¨ker & Ru¨diger (1999). As the brown dwarfs get older, their atmospheres cool and their activity decreases, and it finally ceases at an effective temperature around 2400 K (spectral type M9–L0). The decrease in Fig. 6.—Photon arrival times and energies for the two X-ray–detected the coronal activity level is apparently accompanied by a embedded objects. No evidence for temporal variability is seen. The mean similar decrease of the coronal temperatures: the young photon energy is 4 keV in both cases. BDs in IC 348 and in Oph have coronal temperatures of about 1–2 keV; the quiescent X-ray emission of the M8 field ergs s1 for 2MASS 34443.3+320131. These values are . star VB 10 suggests a temperature of kT 0.5 keV (see slightly lower but still at a comparable level to the X-ray Fleming & Giampapa 2002); for the M9 brown dwarf luminosities of the protostars listed in Feigelson & Mont- LP944-20 a plasma temperature of only 0.26 keV was deter- merle (1999). mined during a flare (Rutledge et al. 2000), while no evi- We also note that L-51 and 2MASS 34443.3+320131 are dence for quiescent X-ray emission from this object has the only two known Class I objects (candidates) in IC 348; been found so far (Martin & Bony 2001). no further similarly red objects have been found up to now in infrared observations (see the near-infrared color-color diagram in Fig. 3 in PZ01). Therefore, the X-ray detection 8. X-RAY PROPERTIES OF THE EMBEDDED OBJECTS frequency among the known Class I objects in IC 348 is In our Chandra image of IC 348 we also discovered X-ray 100%. Of course, this does not exclude the possible existence emission from two deeply embedded objects, presumably of further, more deeply embedded objects. In fact, at least Class I protostars, south of the cluster center. The first one such object is known to exist in the molecular cloud object is L-91, which has been classified as an infrared Class south of IC 348, i.e., the source of the very young molecular I source by L98. The second object is 2MASS hydrogen jet HH 211 (see McCaughrean, Rayner, & Zin- 034443.3+320131. The plots of photon arrival times versus necker 1994) for which the nondetection in our deep K-band photon energies for these two objects are shown in Figure 6. images implies an extinction of at least AV > 50 mag. The photon arrival times show no evidence for significant Besides the unfortunate fact that it lies just outside the field variability in these two objects. The energies of the detected of view of our Chandra image, it would hardly be detectable photons are rather high, with a mean of around 4 keV. This if it were to emit X-rays at similar levels to the two detected is not surprising, given the strong extinction toward these embedded objects. objects, which efficiently absorbs the softer photons. In order to get an estimate of the plasma temperatures, we extracted X-ray spectra and fitted them with a single-tem- 9. RELATION OF THE X-RAY EMISSION TO BASIC perature plasma model plus absorption. As the numbers of STELLAR PARAMETERS detected photons are rather small, the fitting parameters are necessarily associated with rather larger uncertainties. For 9.1. X-Ray Luminosity and Bolometric Luminosity L-51, the fit yields a plasma temperature of kT 7 keV and In order to estimate bolometric luminosities for as many 22 2 a hydrogen column density of NH 8.6 10 cm ; for objects as possible in a homogeneous way, we used the J- 2MASS 34443.3+320131 we find kT 6keVand band magnitudes listed in the 2MASS catalog. The J-band 22 2 NH 4.8 10 cm . These plasma temperatures are flux is a good tracer of the stellar luminosity for late-type clearly higher than those found for the vast majority of T stars because the J band coincides well with the peak of the Tauri stars (see Fig. 2) but in the typical range of plasma spectral energy distribution for cool stars. Also, the J band temperatures found for extremely young embedded objects is a good compromise between shorter wavelength bands, (kT 5–8 keV; see Feigelson & Montmerle 1999; Imanishi which would be too strongly affected by extinction, and lon- et al. 2001a). The X-ray luminosities derived from these fits ger wavelength bands, which would be contaminated by 30 1 30 are LX 1.5 10 ergs s for L-51 and LX 0.9 10 infrared excesses of the T Tauri stars. Using the extinctions listed in Tables 1 and 2, we computed the dereddened J- band magnitude J . Then we employed the most recent ver- 4 Fleming & Giampapa (2002) estimated that the ionization fraction in 0 the upper photosphere drops by about 2 orders of magnitude from sion of the Baraffe et al. (1998) pre–main-sequence (PMS) 5 7 10 for early-M–type objects (Teff & 3300 K) down to 10 for models to construct a relation between J0 and the bolomet- M8 objects (Teff 2600 K). ric luminosity for 1.3 Myr old objects (the average age of the 1622 PREIBISCH & ZINNECKER Vol. 123 stars in IC 348). This relation was then used to determine tion, several two-sample tests, correlation tests, and linear the bolometric luminosities of the objects in IC 348. regressions. The best linear regression fit found with The plot of the extinction corrected 0.2–10 keV band X- ASURV (Isobe, Feigelson, & Nelson 1986) between LX and ray luminosity versus the bolometric luminosity is shown in Lbol (excluding the A- and F-type stars, which are not Figure 7. Nearly all young stars in IC 348 show LX/ coronal X-ray emitters, as well as those sources that showed 5 1 Lbol >10 and therefore are much more X-ray active than X-ray flares) is log LX(ergs s ) = 29.86 + 1.03 6 the Sun (for which LX/Lbol 10 is a typical value). The (0.06) log (Lbol/L) with a scatter of 0.42 in log LX. This 4 most active stars show fractional X-ray luminosities around relation is fully consistent with LX/Lbol =2 10 , within 3 LX/Lbol =10 , which is the saturation limit for coronally the uncertainties and is very similar to the relations found active stars (Fleming, Schmitt, & Giampapa 1995). for other young clusters (see Feigelson & Montmerle 1999). For objects with Lbol & 5 L there is a remarkable bifur- cation of X-ray luminosities: some of these luminous stars 9.2. X-Ray Luminosity and Stellar Mass 30 31 have quite high X-ray luminosities, between 10 and 10 We also considered the relation between stellar mass and 1 ergs s , while for others we find much lower X-ray luminos- X-ray luminosity. For this, we estimated masses5 by com- 29 1 ities around 10 ergs s or very low upper limits at LX/ paring the locations of the objects in the H-R diagram to the . 5 Lbol 10 . These objects are young A- and F-type stars PMS evolutionary models of D’Antona & Mazzitelli (1997). (or very young G-type pre–main-sequence stars that will Figure 8 shows a clear correlation between stellar mass and reach the main sequence as A-type stars), for which no coro- X-ray luminosity. The ASURV linear regression fits gave6 nal emission is expected (e.g., Schmitt et al. 1993). The log L (ergs s1) = 30.10 + 1.97(0.24) log (M/M ), with observed X-ray emission probably originates from unre- X a scatter of 0.80 in log LX. This slope is lower than that solved late-type companions. This presumption is sup- found for the T Tauri stars in the Chamaeleon star-forming ported by the X-ray properties of the two X-ray bright A- region (slope = 3.6 0.6; Feigelson et al. 1993) and also type stars: The A4 star associated with CXOPZ- slightly lower than that derived for M-type field stars J034450.7+321904 showed a strong flare (see x 6) with prop- (slope 2.5; Fleming et al. 1988). The slope agrees with the erties typical for flaring coronal sources. The A2 star associ- slope of the power-law relation between Lbol and stellar ated with CXOPZ-J034435.3+321005 has an X-ray mass (Burrows et al. 2001) and therefore is consistent with spectrum with kT = 2.46 keV, as typically found for active the LX / Lbol relation found above. As the scatter in the LX coronal sources. versus M relation is nearly twice as large as the scatter in the We utilized the ASURV survival analysis package (LaV- LX versus Lbol relation, the X-ray luminosity is better corre- alley, Isobe, & Feigelson 1992) for the statistical investiga- lated with the bolometric luminosity than with the mass. tion of the relation between LX and Lbol. The ASURV software allows one to deal with data sets that contain non- 9.3. Coronal and Chromospheric Emission detections (upper limits), as well as detections, and provides the maximum-likelihood estimator of the censored distribu- The deep optical observations by H98 and L98 have yielded H emission-line equivalent widths for many objects in IC 348. We can use these data to explore the relation between the H emission (as a tracer of the chromospheric activity) and the X-ray emission (as a tracer of the coronal activity). We computed the H luminosities from the R-band magnitudes and H equivalent widths with the method described in Reid, Hawley, & Mateo (1995). A plot of X-ray luminosity versus H luminosity is shown in Figure 9. For the following analysis, we will distinguish7 between the weak-line T Tauri stars [WTTSs; W(H) < 10 A˚ ], in which the H emission can be assumed to be mainly of chromospheric origin, and the classical T Tauri stars [CTTSs, W(H) 10 A˚ ], in which the H emission is probably dominated by the accretion of circumstellar material onto

5 We are aware that these mass estimates are subject to significant uncertainties, since masses estimated with different PMS models and/or temperature scales can disagree by as much as a factor of 2 (a detailed investigation of the uncertainties of mass estimates can be found in L99). However, here we are not so much interested in exact masses of individual objects, as in the general relation between stellar mass and the X-ray luminosity; for this, these uncertainties are not very relevant as long as the rela- tive ordering in mass is correct. Fig. 7.—X-ray versus bolometric luminosity for the objects in IC 348. 6 For the regression fit we excluded the A- and F-type stars, which are Solid dots show the X-ray detected cluster members, arrows indicate upper not coronal X-ray emitters. limits for the undetected cluster members. Sources that showed a large X- 7 We note that a classification scheme slightly different from what we use ray flare during the observation (see x 6) are marked by triangles. The dot- here has been suggested by Martin (1997): it uses a spectral-type–dependent 5 4 ˚ ˚ ted lines show the relations LX/Lbol =10 (bottom line), 10 (middle line), boundary, which is 5 A for K-type stars, 10 A for early M types, and 20 and 103 (top line). The thick dashed line shows the regression fit computed A˚ for late M types. Application of this scheme would change the classifica- with the ASURV survival analysis package, which is consistent with LX/ tion of only two WTTSs to CTTSs and of only three CTTSs to WTTSs. 4 Lbol 2 10 . This would not cause any noticeable difference in our results. No. 3, 2002 X-RAY EMISSION IN IC 348 1623

Fig. 8.—X-ray luminosity vs. stellar mass for the objects in IC 348. The Fig. 9.—X-ray luminosity vs. H luminosity for the T Tauri stars in IC thick dashed line shows the regression fit computed with ASURV, roughly 348. The crosses show the WTTSs [W(H) < 10 A˚ ], the gray dots show the 2 corresponding to LX / M . CTTSs [W(H) 10 A˚ ]. The arrows mark upper limits for the X-ray luminosities of undetected objects. The two dotted lines show the relations LX/ LH = 1 and LX/LH = 1 0. The thick dashed line shows the regression fit. the young stars. For the WTTS, Figure 9 shows a very good correlation between LX and LH. A linear regression fit with 1 ASURV yields log LX(ergs s ) = 30.11 + 0.83(0.11) 30 1 (Fig. 10) and those of the infrared-excess sources versus the log [LH/(10 ergs s )]. This is consistent with the slope of nonexcess sources (Fig. 12). 0.7 between log LX and log LH reported by Fleming et al. (1988) for their sample of coronally active nearby M dwarfs. The CTTSs, on the other hand, strongly deviate from this 10.1. Classical versus Weak-Line T Tauri Stars relation, which is no surprise as their strong H emission Many X-ray observations of star-forming regions have comes mainly from accretion processes. been used to search for differences in the X-ray properties of Nearly all of our WTTSs show fractional H luminosities CTTSs [W(H) 10 A˚ ] as compared with WTTSs between L /L 104.5 and L /L 103.5, just as H bol H bol [W(H) < 10 A˚ ]. The results of these investigations show found for low-mass Hyades and Pleiades dwarfs (Reid et al. remarkable differences between individual regions. In most 1995) and nearby low-mass stars (Gizis et al. 2000). Also, regions, e.g., the Orion nebula cluster (Gagne´ et al. 1995), the ratio of X-ray to H luminosity for our WTTSs shows a the Chameleon I cloud (Feigelson et al. 1993), and the very similar distribution between L /L 1 and L / X H X Ophiuchi cloud core (Casanova et al. 1995), no significant L 10 as the ratios observed for the low-mass Hyades H differences could be found between the X-ray luminosity and Pleiades dwarfs (Reid et al. 1995). This suggests that the functions of CTTSs and WTTSs. In the Taurus-Auriga coronal and chromospheric properties of the WTTSs in IC region, however, the arguments for strong differences in the 348 are very similar to those of late-type main-sequence X-ray luminosity functions of CTTSs and WTTSs reported stars. This supports the idea that coronal and chromo- by Neuhaüser et al. (1995) have not only been confirmed spheric properties of the very young WTTSs can be under- but even strengthened: Stelzer & Neuhaüser (2001) studied stood as due to scaled up solar-type dynamo activity. Like the X-ray properties of 170 young stars in Taurus-Auriga Fleming et al. (1988), we conclude that these results support and found that the X-ray luminosity functions of CTTSs the hypothesis that the chromosphere is heated by X-rays and WTTSs are clearly different, with the WTTSs being the from the overlying corona. stronger X-ray emitters; the median X-ray luminosity of the K- and M-type WTTSs is a factor of 6 higher than that of the CTTSs, if upper limits for the nondetected objects are 10. X-RAY EMISSION AND CIRCUMSTELLAR included. As explained below, we suspect that this result is PROPERTIES due to selection effects. Up to now it is not clear whether the X-ray activity of T The combination of our deep Chandra data with very sen- Tauri stars, especially the CTTSs, can be fully explained by sitive optical surveys (H98, L98) makes IC 348 an ideal tar- scaled up solar-type dynamo activity, or whether nonsolar get to study this question. In Figure 10 we plot the X-ray magnetic structures, which couple the star to its circumstel- luminosities of the stars as a function of W(H) and LX ver- lar disk, might possibly be involved (see Shu et al. 1997; sus Lbol for WTTSs and CTTSs. Note that these compari- Walter & Byrne 1998). Despite the abundance of observa- sons are based on a purely H-selected sample of T Tauri tional results, this question has not yet been fully answered. stars, which is free of any X-ray–selection bias. Considering To gain more insight into this problem, we will compare the only the X-ray–detected T Tauri stars, there is no significant X-ray properties of classical versus weak-line T Tauri stars difference in the X-ray properties of the WTTSs as com- 1624 PREIBISCH & ZINNECKER Vol. 123

Fig. 10.—Top left: X-ray luminosities and upper limits for all optically selected T Tauri stars in IC 348 vs. their H emission-line equivalent widths. Top right: X-ray luminosity vs. bolometric luminosity for WTTSs (crosses) and CTTSs (dots). Upper limits are marked by arrows. One can see that, due to the large number of nondetections among the faint CTTSs, the median X-ray luminosity of the WTTSs is higher than that of the CTTSs. This, however, has nothing to do with the presence or absence of circumstellar disks around the T Tauri stars (see text). Bottom: Kaplan-Meier estimator for the X-ray luminosity functions of the H selected CTTSs and WTTSs, which are clearly diﬀerent.

pared with the CTTSs. However, many CTTSs were not 1. H emission is thought to trace the current accretion detected as X-ray sources; including these nondetections, activity, rather than the presence of circumstellar material; the conclusion changes. To obtain a quantitative assess- the accretion activity in T Tauri stars is thought to be highly ment, we used the ASURV package to construct the variable in time (see Calvet, Hartmann, & Strom 2000). Kaplan-Meier estimator for the X-ray luminosity function 2. H emission is known to be strongly time variable in (the cumulative distribution function [CDF]) of the H- many T Tauri stars (e.g., Guenther & Emmerson 1997); selected CTTSs and WTTSs. The two distribution func- objects with variable H emission may appear as WTTSs at tions, also shown in Figure 10, are clearly different, and the one time and as CTTSs at another time. two sample tests implemented in ASURV give probabilities 3. M-type stars without any circumstellar material can between P = 0.003 and P = 0.07 for the null hypothesis show quite strong chromospheric H emission (see x 9.3); it that both distributions are equal. This suggests a significant is not straightforward to distinguish between the chromo- difference in the X-ray properties of the WTTSs as com- spheric H emission and the H emission from accretion. pared with the CTTSs, similar to that found for the Taurus- 4. Thus the use of H emission as a tracer to find T Tauri Auriga stars mentioned above. stars can introduce very serious selection effects when it However, we would like to point out that there are major comes to WTTSs. CTTSs are very easy to recognize by their problems related to the use of H emission as an indicator prominent H emission, even if they are very faint (such as, for circumstellar material: for example, the faint LH T Tauri stars in IC 348), whereas No. 3, 2002 X-RAY EMISSION IN IC 348 1625

WTTSs of similar brightness are much harder to identify; in fact, the majority of WTTSs in many star-forming regions has been found in X-ray observations (see Neuha¨user 1997). Therefore, if we wish to obtain a complete, brightness- limited sample of T Tauri stars (both CTTSs and WTTSs), it is necessary to use photometric or spectroscopic selection criteria. Indeed, this was done for IC 348, allowing us to proceed further and study the X-ray properties of young stellar objects with and without circumstellar material (proxy for CTTSs and WTTSs, respectively).

10.2. L-Band Excess Sources versus Nonexcess Sources A very good way to distinguish between young stellar objects with and without circumstellar material is to look for near-infrared excess emission. While JHK photometry alone is not sufficient to detect circumstellar material, Haisch et al. (2001b) have demonstrated that JHKL photometry allows an almost unambiguous discrimination between stars with and without circumstellar material. For the stars in IC 348, we can take advantage of the recent photometric survey reported by Haisch, Lada, & Lada (2001a), Fig. 11.—Infrared excess D(KL) vs. H emission-line equivalent width which provided JHKL photometry for 107 sources and for the stars in IC 348. The solid dots show the X-ray detected objects, the ˚ found 65% of the stars to have significant infrared excess. crosses the nondetected objects. The vertical line at W(H)=10Ais the boundary between WTTSs and CTTSs; the horizontal line at We used these data to determine the infrared excess from D(KL) = 0.17 is the boundary between infrared excess sources (Ex) and the JHKL colors of the stars in the following way: We nonexcess sources (noEx). It can be seen that not only the CTTSs, but even placed the objects in the (KL) versus (JH ) color-color many WTTSs have infrared excess (circumstellar material). diagram (see Fig. 7 in PZ01) and determined the quantity D(KL), which we defined as the horizontal displacement of each object from the blueward boundary of the photo- sample. The JHKL photometric sample of Haisch et al. spheric reddening band. The quantity D(KL) is therefore a (2001a) is limited by the L-band magnitude of the stars. This measure of the infrared excess emission corrected for the clearly introduces a selection effect for stars with strong reddening due to extinction. Given the spread in intrinsic infrared excess, because these will be brighter in the L band photospheric (KL) colors for late-type stars, values of and therefore easier to detect. Our aim therefore was to D(KL) > 0.17 can be considered evidence for infrared select a subsample of objects that is complete in the sense excess; stars with lower D(KL) may either have a weak that all stars of a given luminosity are included, irrespective infrared excess or may be intrinsically very red (i.e., very late of their infrared excess. As the J-band flux is a good proxy type) objects. We therefore classified all stars with for the stellar luminosity of cool late-type stars, we used the D(KL) > 0.17 as excess sources (Ex), and those with J-band magnitude as a selection criterion. The sample of D(KL) 0.17 as nonexcess sources (noEx). Haisch et al. (2001a) appears to be complete for objects with In Figure 11 we compare the infrared excesses of the stars J 12.5; below that limit, there is a selection effect favoring in IC 348 with their H emission-line equivalent widths. All strong excess sources. CTTSs (but one) show strong infrared excess. This is not In Figure 12 we plot the X-ray luminosities of the stars in very surprising, because, if H emission traces accretion IC 348 against their infrared excesses and X-ray versus activity, it is necessarily tied to the presence of circumstellar bolometric luminosity for excess sources and nonexcess material that causes the infrared excess. A more interesting sources. The large symbols mark the objects in the bright- result is obtained for the WTTSs: the majority of the ness-limited ‘‘ complete sample ’’; fainter objects are shown WTTSs shows infrared excess, in several cases of compara- by small symbols. These plots show no systematic differen- ble strength to the CTTSs. This clearly shows that many of ces in the X-ray luminosities of excess sources versus non- the WTTSs are actually surrounded by significant amounts of excess sources. The two-sample tests in ASURV gave prob- hot circumstellar material. When looking only at their H abilities between P =0.54andP = 0.96 that the distribu- emission, one would misclassify these objects as ‘‘ naked,’’ tions of X-ray luminosities are equal for both samples. whereas in fact they do have circumstellar material but cur- In conclusion, we find no evidence that the presence of rently are accretion quiet. This result demonstrates that the circumstellar material makes any difference to the X-ray infrared excess gives a more realistic picture of the circum- properties of the T Tauri stars. The apparent difference in stellar properties of the T Tauri stars than the H emission the X-ray luminosity functions of WTTSs versus CTTSs is (see Haiseh et al. [2001b]). A similar conclusion was drawn probably caused by the fact that H emission is a poor indi- by Andre´ & Montmerle (1994) based on a sample of Tau- cator of circumstellar material around T Tauri stars and by rus-Auriga T Tauri stars. These authors suggested that the strong selection effects favoring the optical detection of faint WTTSs with infrared excess might be in a transition phase, CTTSs (which are faint in X-ray too) with strong H emis- becoming naked T Tauri stars by inside-out clearing of their sion, as opposed to faint WTTSs. The lack of faint WTTSs disks. in the sample of young stars tends to shift the median X-ray An important advantage of the infrared excess classifica- luminosity of WTTSs to higher values, while the inclusion tion is that it allows one to consider the completeness of the of faint CTTSs decreases the median X-ray luminosity of 1626 PREIBISCH & ZINNECKER Vol. 123

Fig. 12.—Top left: X-ray luminosity of the T Tauri stars in IC 348 vs. their near-infrared excess. The T Tauri stars in the complete sample are shown by the big dots, the fainter T Tauri stars, for which the sample is incomplete, by the small crosses. Top right: X-ray luminosity vs. bolometric luminosity for nonexcess sources (crosses) and the excess sources (dots). Upper limits are marked by arrows. Bottom: Kaplan-Meier estimator for the X-ray luminosity functions of the nonexcess sources and the excess sources. Note that there is no signiﬁcant diﬀerence in these distributions. These plots should be compared with the corresponding ones in Fig. 10, although the sets of stars are not identical.

CTTSs. In our view, this explains the apparent difference in peratures of the T Tauri stars and their infrared excesses or the X-ray luminosity functions between the CTTSs and their H emission-line widths. WTTSs in IC 348.8 We also mention in passing that we Finally, we note that our result agrees very well with a could not find any relation between the X-ray plasma tem- similar investigation for the T Tauri stars in the Oph star-forming region by Grosso et al. (2000), who also found no statistically significant differences between the X- 8 We believe that the Taurus results are caused by similar selection ray luminosity functions of T Tauri stars with and without effects. In this context it is interesting to consider the numbers of CTTSs disks. Our results provide no evidence for any nonsolar- and WTTSs among different spectral types in the T Tauri star sample used like coronal contribution to the X-ray emission of the T by Stelzer & Neuhaüser (2001): there are many more (36) WTTSs than Tauri stars. If, for example, a disk dynamo were at work, CTTSs (22) among the K-type stars, while among the M-type stars there are many more (61) CTTSs than WTTSs (34). These numbers seem to one would wrongly predict systematically higher X-ray reflect the observational selection effect for the (rather easy to find) CTTSs luminosities for the CTTSs (or excess sources). However, among the fainter (M-type) stars. Furthermore, we note that in a large we note that these considerations are valid only for the T extended association such as Taurus, it is extremely hard (perhaps impossi- Tauri stars, which have ages of &1 Myr; the X-ray emis- ble) to get an optically complete sample of all young stars. Only in rather compact clusters such as IC 348 (or in small subregions of associations; see, sion from much younger (embedded) objects might well be e.g., Preibisch, Guenther, & Zinnecker 2001) is it possible to obtain com- (partly) caused by other mechanisms (see Feigelson & plete samples of young stars by optical/infrared spectroscopic surveys. Montmerle 1999). No. 3, 2002 X-RAY EMISSION IN IC 348 1627

11. SUMMARY is a tracer of chromospheric activity. This suggests that the The main results of our analysis of the Chandra ACIS coronal and the chromospheric activity of the WTTSs has a observation of IC 348 can be summarized as follows: common origin in an enhanced solar-like dynamo effect. 7. The classical T Tauri stars [W(H) 10 A˚ ] appear to 1. The analysis of the X-ray spectra yields plasma tem- be systematically weaker X-ray emitters than the weak-line peratures between 0.7 and 3 keV for most sources and T Tauri stars [W(H) < 10 A˚ ]. However, this apparent dif- higher temperatures, up to 6 keV, for some sources. This is ference is most likely caused by a selection effect related to the typical range for coronally active stars. the H emission (faint WTTSs are much harder to find than 2. About half of the X-ray sources in IC 348 show tempo- faint CTTSs). If we use the presence or absence of near- ral variability; some sources show large X-ray flares. Time- infrared L-band excesses to distinguish between objects with resolved spectroscopy of the strongest flares yields plasma and without circumstellar material, we find no differences in temperatures around 6 keV during the flare peak. In several the X-ray properties of the two samples. This strongly sug- flare light curves the peak in the hard X-ray (2–10 keV) gests that there is no causal relationship between the circum- emission occurs about 1 hr earlier than the peak in the soft stellar environment and the X-ray properties of the T Tauri X-ray (0.2–2 keV) emission. stars. 3. The X-ray emission detected from seven very low mass 8. The X-ray properties of young low-mass stars, even and substellar objects in IC 348 suggests plasma tempera- those with circumstellar disks, suggest that the emission is tures around 1–2 keV. Three objects are probably varia- generated by a stellar dynamo. ble; one is clearly variable and perhaps showed a moderate X-ray flare during our observation. These properties suggest that the X-ray emission from the brown dwarfs originates from a hot corona, as commonly observed for late-type T. P. would like to thank George Herbig and the other stars. Since the very young substellar objects in IC 348 are members of the Institute for Astronomy at the University of still of spectral type M, it is not surprising that their coronal Hawaii for their kind hospitality during his visit in 2001 properties are similar to late-M–type stars, the more so, February, where part of this paper was written, and for since both types of objects are fully convective (same interesting and useful discussions on IC 348. H. Z. would dynamo process). like to thank Gu¨nther Hasinger, director of the Astrophysi- 4. The two X-ray–detected optically invisible young stel- kalisches Institut Potsdam, for his support over the last few lar objects south of IC 348 emit hard X-rays with a tempera- years. We are grateful to the referee, T. Montmerle, for his ture of 6 keV, as is typical for the X-ray emission from very constructive referee’s report, which helped to improve infrared Class I protostellar objects. We find no evidence for this paper. We acknowledge the use of the ASURV survival X-ray flares on these objects. analysis package (LaValley et al. 1992) which is freely avail- 5. We find a strong correlation between the X-ray lumi- able from http://www.astro.psu.edu/statcodes/asurv. This nosity and bolometric luminosity of the stars in IC 348, con- publication makes use of data products from the Two 4 sistent with LX/Lbol 2 10 . We also find a good Micron All Sky Survey, which is a joint project of the Uni- correlation between X-ray luminosity and stellar mass versity of Massachusetts and the Infrared Processing and 2 (LX / M ). Analysis Center, funded by the National Aeronautics 6. For the WTTSs in IC 348 we find a good correlation and Space Administration and the National Science between the X-ray luminosity and the H luminosity, which Foundation.

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