THE MASS ACCRETION RATES OF INTERMEDIATE-MASS T TAURI STARS

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1 The Astronomical Journal, 128: , 2004 September # The American Astronomical Society. All rights reserved. Printed in U.S.A. THE MASS ACCRETION RATES OF INTERMEDIATE-MASS T TAURI STARS Nuria Calvet, 1, 2 James Muzerolle, 3, 4 César Briceño, 2, 5 Jesus Hernández, 2, 5 Lee Hartmann, 1 José Luis Saucedo, 1, 6 and Karl D. Gordon 3 Received 2004 January 14; accepted 2004 May 19 ABSTRACT We present Hubble Space Telescope ultraviolet spectra and supporting ground-based data for a sample of nine intermediate-mass T Tauri stars (IMTTSs; M ). The targets belong to three star-forming regions: T Tau, SU Aur, and RY Tau in the Taurus clouds; EZ Ori, P2441, and V1044 Ori in the Ori OB1c association surrounding the Orion Nebula cluster; and CO Ori, GW Ori, and GX Ori in the ring around k Ori. The supporting groundbased observations include nearly simultaneous UBV(RI) C photometry, 6 8 resolution spectra covering the range , optical echelle observations in the range , andk-band near-infrared spectra. We use these data to determine improved spectral types and reddening corrections and to obtain physical parameters of the targets. We find that an extinction law with a weak feature but high values of A UV =A V is required to explain the simultaneous optical-uv data; the reddening laws for two B-type stars located behind the Taurus clouds, HD and HD , meet these properties. We argue that reddening laws with these characteristics may well be representative of cold, dense molecular clouds. Spectral energy distributions and emission-line profiles of the IMTTSs are consistent with expectations from magnetospheric accretion models. We compare our simultaneous optical-uv data with predictions from accretion shock models to get accretion luminosities and mass accretion rates (Ṁ ) for the targets. We find that the average mass accretion rate for IMTTSs is 3 ; 10 8 M yr 1,afactorof5 higher than that for their low-mass counterparts. The new data extend the correlation between Ṁ and stellar mass to the intermediate-mass range. Since the IMTTSs are evolutionary descendants of the Herbig Ae/Be stars, our results put limits to the mass accretion rates of their disks. We present luminosities of the UV lines of highly ionized metals and show that they are well above the saturation limit for magnetically active cool stars but correlate strongly with accretion luminosity, indicating that they are powered by accretion, in agreement with previous claims but using a sample in which reddening and accretion luminosities have been determined self-consistently. Finally, we find that the relation between accretion luminosity and Br luminosity found for low-mass T Tauri stars extends to the intermediate-mass regime. Key words: accretion, accretion disks circumstellar matter stars: formation stars: pre main-sequence 1. INTRODUCTION Many efforts have been made in recent years to quantify the rate of mass accretion in disks around young stellar objects (YSOs). The mass accretion rate Ṁ is a key parameter for understanding disk structure and evolution, as well as planet formation and migration. Using detailed models of disk structure, Ṁ together with the viscosity or rate of angular momentum transport determines the radial dependence of the disk surface density, providing predictions that can be compared with observations (D Alessio et al. 1999, 2001). Mass accretion rates generally appear to decrease with time as mass is accreted onto the star and disks grow to conserve angular momentum (Hartmann et al. 1998; Hartmann 1998), making Ṁ an indicator 1 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138; ncalvet@cfa.harvard.edu, hartmann@cfa.harvard.edu. 2 Also Postgrado de Fisica Fundamental, Universidad de Los Andes, Mérida, Venezuela. 3 Steward Observatory, 933 North Cherry Avenue, University of Arizona, Tucson, AZ Visiting Astronomer, Kitt Peak National Observatory, National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation. 5 Centro de Investigaciones de Astronomía, Apdo. Postal 264, Mérida 5010-A, Venezuela. 6 Instituto de Astronomía, Universidad Nacional Autónoma de México, Apdo. Postal , México, DF, Mexico of the evolutionary stage of the disk. Moreover, giant planet formation is critically dependent on the disk surface density (Pollack et al. 1996). Early attempts to measure mass accretion rates from disk emission were invalidated when it was recognized that stellar irradiation is the most important mechanism of disk heating in most systems (Kenyon & Hartmann 1987). The best determination of Ṁ comes from measuring the luminosity of the excess emission that veils the intrinsic photospheric spectrum of a YSO. In the presently accepted paradigm, this excess arises from the accretion shock formed at the stellar surface as disk material falls along stellar magnetic field lines, which truncate the disk at a few stellar radii (Uchida & Shibata 1985; Bertout et al. 1988; Königl 1991; Camenzind 1990; Shu et al. 1994). The magnetospheric mode of accretion for low-mass T Tauri stars (<1 M ) is supported by many lines of evidence, among them the characteristic broad line profiles consistent with formation in the extended magnetospheric flow (Muzerolle et al. 1998a, 1998b, 2001), the strength and geometry of the magnetic field capable of truncating the inner disk (Johns-Krull & Hatzes 1997; Johns-Krull et al. 1999a, 1999b), and the spectral energy distribution of the optical and ultraviolet (UV) excess consistent with accretion shock emission (Calvet & Gullbring 1998, hereafter CG98; Gullbring et al. 2000; Ardila et al. 2002, hereafter ABWVJ02). Measurements of the luminosity of the excess flux above the photosphere in the low-mass classical T Tauri stars (CTTSs; 0:1 M M 1 M ) have yielded mass accretion

2 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1295 TABLE 1 Stellar Properties Object Region Spectral Type T eff (K) L (L ) R (R ) M Age (M ) A V ( Myr) T Tau... Taurus G SU Aur... Taurus G RY Tau... Taurus G EZ Ori... Ori OB 1c G P Ori OB 1c F V1044 Ori... Ori OB 1c G CO Ori... k Ori G GW Ori... k Ori G GX Ori... k Ori G Note. Errors in spectral type determinations are plus or minus two subclasses. Distances are taken as 140 pc for Taurus ( Kenyon et al. 1994), 440 pc for Ori OB1c (Genzel et al. 1981), and 450 pc for the k Ori ring ( Dolan & Mathieu 2001). rates Ṁ for these stars (Hartigan et al. 1991, 1995; Valenti et al. 1993; Gullbring et al. 1998). Using calibrations of accretion luminosity with either excess luminosity based on broadband photometry or emission-line luminosity, the sample of CTTSs with known Ṁ is presently very large (Hartmann et al. 1998; White & Ghez 2001; White & Basri 2003). In addition, magnetospheric modeling of emission-line profiles in very low mass young objects has now yielded Ṁ for stars down to the substellar limit (Muzerolle et al. 2001, 2003a, hereafter M03a). The measurement of Ṁ has proved much more difficult for YSOs of higher masses. The best-known intermediate-mass YSOs are the Herbig Ae/Be stars (HAeBe; Herbig 1960), and Muzerolle et al. (2003b) have shown that at least for many late B to late A stars of the class (HAe; Natta et al. 2001), emission-line profiles are consistent with magnetospheric accretion. However, since the effective temperature of these stars is similar to the temperatures characterizing the accretion shock emission in the optical and near-uv (NUV), it is difficult to effectively isolate the excess emission and measure its luminosity (Muzerolle et al. 2003b). The intermediate-mass T Tauri stars (IMTTSs; also called ETTSs by Herbst et al and GTTSs by Herbst & Shevchenko 1999) are a subset of the T Tauri class with masses 1 M M 5 M, the same mass range as the HAe. They are already on radiative pre main-sequence (PMS) evolutionary tracks but have spectral types ranging from early K to late F. These stars are the evolutionary predecessors of the Herbig Ae stars, so an understanding of disk evolution in intermediate-mass stars must incorporate both the IMTTSs as well as the HAe stars. There are few measurements of Ṁ in IMTTSs. For one thing, they are less numerous than low-mass TTSs; in addition, their photospheres are much brighter than TTSs in the blue, in which deveiling procedures to extract excess spectra are usually done. More accurate measurements of the excess can be done in the UV, in which the separation between the photospheric emission, 5000 K < T < 6200 K, and the shock emission, T shock 8000 K, is clearer. Gullbring et al. (2000) used IUE spectra for three IMTTSs to attempt a first measurement of Ṁ. However, the wavelength coverage and signal-to-noise ratio of the IUE spectra were limited. In this paper, we present and analyze Hubble Space Telescope (HST ) Space Telescope Imaging Spectrograph (STIS) UVspectraforasampleofnineIMTTSsandmeasureaccretion luminosities and mass accretion rates from the excess emission over photospheric fluxes. We also present a number of supporting observations to the HST measurements to improve the reliability and extend the scope of the accretion determinations. We secured nearly simultaneous optical observations for most of our sample, which helped us fix photospheric flux levels. We obtained large-wavelength coverage spectra to redetermine in a consistent fashion spectral types of the observing sample. Parallel to this effort, we obtained optical echelle spectra, which allowed us to further assess the magnetospheric model of accretion. Finally, we obtained nearinfrared spectra, which allowed us to extend the calibration between Br luminosities and accretion luminosities over 1 order of magnitude in mass. Overall, we present a study of a subgroup of TTSs, using for the first time simultaneous optical and UV data, which allows us to determine reddening and stellar and accretion properties self-consistently, with applications for less well-studied samples. 2. OBSERVATIONS 2.1. HST Observvations Our sample was selected among objects classified as late F to early K in the catalog of Herbig & Bell (1988), such that they were bright and had relatively low reddening estimates. With these conditions, we selected nine objects located in three star-forming regions: the Taurus-Auriga clouds, the subassociation Ori OB1c, which encompasses the Orion Nebula Cluster (ONC) region, and the ring around k Ori. The targets and their locations are shown in Table 1, together with stellar properties that are discussed in x 3.3. The sample was observed with HST in program GO8317, using the STIS MAMA detector, with G230L and G140L gratings to cover the NUV and far-uv (FUV) regions, respectively. The resolving power of these gratings is 1.58 and pixel 1, respectively; with FWHM of 1.5 and 2 pixels, this resulted in a resolution of 3 8 in the NUV and 1 8 in the FUV. A 2 00 ; slit was employed, and the data were taken in binned pixels mode on dates and exposure times indicated in Table 2. For one of the targets, GX Ori, we only obtained a NUV spectra, because it was estimated to be too faint to reach the FUV in reasonable exposure times. Standard IRAF 7 STSDAS/HST_CALIB/STIS/CALSTIS pipeline procedures 7 IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

3 1296 CALVET ET AL. Vol. 128 TABLE 2 Log of HST Observations Object Instrument-Configuration Start Time (UT) Exposure (s) T Tau... STIS/FUV-MAMA 2000 Feb 7 3:04: STIS/ NUV-MAMA 2000 Feb 7 3:38:39 40 SU Aur... STIS/FUV-MAMA 2000 Feb 13 15:16: STIS/ FUV-MAMA 2000 Feb 13 16:42: STIS/ NUV-MAMA 2000 Feb 13 17:27:06 88 RY Tau... STIS/ NUV-MAMA 2000 Sep 20 7:32:2 288 STIS/ FUV-MAMA 2000 Sep 20 7:44: EZ Ori... STIS/ NUV-MAMA 2000 Feb 25 9:06: STIS/ FUV-MAMA 2000 Feb 25 9:17: P STIS/FUV-MAMA 2000 Sep 20 6:00: STIS/ NUV-MAMA 2000 Sep 20 6:34:24 40 V0144 Ori... STIS/ NUV-MAMA 2000 Feb 15 20:19: STIS/ FUV-MAMA 2000 Feb 15 20:31: CO Ori... STIS/FUV-MAMA 2000 Feb 25 15:30: STIS/ FUV-MAMA 2000 Feb 25 16:55: STIS/ NUV-MAMA 2000 Feb 25 17:37: GW Ori... STIS/FUV-MAMA 2000 Feb 22 18:16: STIS/ FUV-MAMA 2000 Feb 22 19:41: STIS/ NUV-MAMA 2000 Feb 22 20:25:3 88 GX Ori... STIS/ NUV-MAMA 2000 Feb 20 17:57: were used to reduce the data from flat-fielded, bias-corrected wavelength-dispersion images, which in turn were processed through the BASIC2D task. The default pipeline in the spectral extraction X1D task automatically corrects the wavelengths to an heliocentric reference frame and performs background extraction from a default or arbitrary sky offset from the star. The signal-to-noise ratio per pixel of the target spectra in the NUV was quite high, 100, except for the case of P2441, with 20. In the FUV, the signal-to-noise ratio was lower, 10 50, except for T Tau and RY Tau. For one of the targets, T Tauri, the FUV spectrum shows extended emission beyond the FWHM of the stellar spectrum in the slit direction, P:A: ¼ 30. Saucedo et al. (2003) analyzed the spectra and the spatial distribution of this extended emission. They found that it is dominated by fluorescent H 2 emission corresponding to a cascade down from the upper level of the 1 2P(5) transition, which in turn is pumped by radiation of Ly at +98 km s 1 from line center. The brightest emission near the star follows the envelope cavity (Stapelfeldt et al. 1998) opened by the stellar outflow. We examined the rest of the sample and found no indication of similar emission extending beyond the width of the stellar spectra in the slit axis direction. We carried out an analysis of the point-spread function (PSF) similar to that in Herczeg et al. (2002) for TW Hydra. We selected small boxes on the wavelength-dispersion images, centered either on a line or the continuum, from which we extracted an average PSF. We find that the wings of the PSF along the slit direction differ from (are higher than) those in the theoretical PSF described on the STIS manual. The discrepancy is the same for every star in the sample, which points toward an instrumental origin rather than real emission. In any event, the amount of extra flux either in lines or continuum is negligible Ground-based Optical Photometry Analysis of the photometric database by Herbst & Shevchenko (1999), which includes six of our targets, indicates that most of them show UX Ori type variability. UX Ori objects suffer occasional obscuration events that can dim the star during several days, even by several magnitudes. One of the targets, CO Ori, is highly variable on timescales of days. For this reason, we considered necessary to obtain groundbased optical photometry on dates as close to the HST observations as possible. Our photometric campaign turned out to be quite successful, and we were able to obtain UBVRI photometry for eight of the targets within a few hours of the HST observations, which is less than the characteristic timescale of UX Ori type variability. Optical photometry was obtained in queue mode with the ANDICAM optical/ir camera on the Yale-Lisbon-Ohio (YALO) 1 m telescope at CTIO (Bailyn et al. 1999), and with the 4-shooter CCD mosaic camera on the SAO 1.2 m telescope at the Fred Lawrence Whipple Observatory (FLWO) on Mount Hopkins, Arizona. In Table 3 we show the observations log, indicating the dates each star was observed. The optical channel of ANDICAM consists of a Loral 2048 ; 2048 CCD with 0B3 pixel 1 scale and a 10A2 ; 10A2 fieldof view. The 4-shooter camera contains four 2048 ; 2048 Loral CCDsseparatedby45 00 andarrangedina2; 2grid.After binning 2 ; 2 during readout, the plate scale was 0B67 pixel 1. In order to achieve more uniform and consistent measurements, we placed all the target standard stars on the same detector (chip 3) of the 4-shooter. The seeing throughout the YALO observations was around 1B5 and during the SAO run was between 1B2 and 2B5. Preliminary reduction of YALO data to make the bias, overscan, and flat-field corrections was handled using customized scripts, developed at Yale University specifically for YALO data, which invoke various IRAF routines. The basic processing of the FLWO data was done with IRAF routines in the standard way. For both runs, the U-band data were flat-fielded using sky flats taken at dusk/dawn. Instrumental magnitudes were obtained using the PHOT package in IRAF, with an aperture of 2.5 times the typical FWHM in the images. The sky background was subtracted by using an annulus with an inner radius of 3 times the stellar FWHM and a typical width of 5 pixels. This

4 No. 3, 2004 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1297 TABLE 3 Log of Optical Observations Object V B V U B V R V I Date Telescope T Tau Feb 8 00:48:15 YALO SU Aur Feb 8 00:20:37 YALO RY Tau Sep 20 11:59:45 48 inch Sep 22 09:52:19 48 inch Sep 22 10:03:16 48 inch EZ Ori Feb 25 01:18:33 YALO Feb 26 01:15:10 YALO P Sep 20 11:12:41 48 inch Sep 20 11:30:06 48 inch Sep 20 11:39:41 48 inch Sep 22 11:21:44 48 inch V1044 Ori Feb 20 01:00:31 YALO CO Ori Feb 25 01:10:12 YALO Feb 26 01:03:46 YALO GW Ori Feb 22 00:35:10 YALO GX Ori Feb 20 00:28:08 YALO approach is appropriate because our target stars are located in sparse regions where crowding is not an issue. The instrumental magnitudes were calibrated in the Johnson UBV and Cousins RI system with observations of Landolt (1992) standard fields containing stars with a range of colors similar to our target objects. Observations of the Landolt standard fields at various air masses could be obtained only for the night of 2000 February 8 at YALO and 2000 February 21 at FLWO. For the remaining nights we derived only zero points and color terms and applied average extinction coefficients for CTIO and FLWO (these differ by only a few hundredths of a magnitude from the values we determined for February 8 and 21). Because our stars are fairly bright sources, measurement errors are negligible, and the main uncertainty is in the transformation to the standard system. Our overall fit errors are 0.06, 0.03, 0.01, 0.02, and 0.02 mag for UBV(RI ) C, respectively. The nightly variation in the zero points was 0.14, 0.06, 0.04, 0.03, and 0.04 mag for UBV(RI ) C, respectively, suggesting reasonable photometric stability throughout the run; these values can also be taken as indicative of the maximum error we can expect in our photometry Spectroscopy Because some of the targets had differing estimates of spectral types in the literature, we obtained spectra to classify them in a consistent system. The spectra were obtained on the 1.5 m telescope of the Whipple Observatory with the FAST spectrograph (Fabricant et al. 1998), equipped with the Loral 512 ; 2688 CCD. The spectrograph was set up in the standard configuration used for FAST COMBO projects, a 300 groove mm 1 grating and a 3 00 wide slit. This combination offers of spectral coverage centered at , with a resolution of 6 8. They were reduced using software developed at SAO specifically for FAST COMBO observations. We also obtained echelle spectra of the targets, which allowed us to obtain detailed line profiles and measure any veiling of photospheric absorption lines. These observations were done with the echelle spectrograph on the 4 m telescope at KPNO on 1999 September Our spectral coverage was , with a resolution R 30; 000. All the spectra were reduced using the IRAF echelle packages, following procedures in Muzerolle et al. (1998b). Finally, we obtained near-infrared spectra with the Cryogenic Spectrometer on the 2.1 m telescope at KPNO on 1998 January We observed Br for each object, with wavelength ranges and m, centered on each line, respectively. The resolution of the spectra is R 800. The spectra were reduced using the standard IRAF routines; details of the data reduction are in Muzerolle et al. (1998c, hereafter MHC98). Br equivalent widths are given in Table ANALYSIS OF THE OBSERVATIONS 3.1. Spectral Types and Veilingg We classified the targets following the spectral classification scheme of Hernández et al. (2004), which is optimized for the wavelength range studied. This scheme is based on 33 spectral features sensitive to changes in T eff, calibrated using piecewise first-order fits. Each index is largely insensitive to reddening, stellar rotation, luminosity class, and signal-to-noise ratio. The spectral type is determined by computing a weighted average of the individual spectral types calculated from each index within a specific spectral-type range. This scheme is designed to avoid indices contaminated by features originating outside the photosphere. In particular, we avoid features like the Na i lines at or the hydrogen Balmer lines. In addition, the large number of indices used allows us to discard those that yield spectral types largely different from the average and that may be contaminated by emission, for example, some of the Object TABLE 4 Br Measurements EW (8) Observed EW (8) K T Tau SU Aur RY Tau EZ Ori P V1044 Ori CO Ori GW Ori GX Ori

5 1298 CALVET ET AL. Vol. 128 Fig. 1. FAST spectra of P2441, CO Ori, GW Ori, SU Aur, and RY Tau. Spectra of a G0 and a G2 standard star are also shown for comparison. The spectral types of the targets fall within this range; see Table 1. Relevant spectral features are identified. Fe ii multiples (see Hernández et al for details). The large wavelength coverage of the FAST spectra, , compensates for the moderate resolution by allowing us to probe features both in the red and the blue regions of the spectra. Optical spectra of the targets and standard stars are shown in Figures 1 and 2, arranged according to the spectral types given in Table 1. The estimated error for the spectral types is plus or minus two subclasses. We indicate the positions of the most important spectral indices used in our classification scheme and some emission lines frequently seen in active PMS stars. As can be seen in Figures 1 and 2, P2441, CO Ori, GW Ori, SU Aur, RY Tau, and V1044 Ori have spectral types similar to or earlier than the standard star G2. However, our spectral types for P2441, CO Ori, GW Ori, and RY Tau differ by more than three subclasses from the later types obtained by Cohen & Kuhi (1979, hereafter CK79), which have been the most frequently used in recent studies. Our spectral type determinations are based mainly on Ca i k4226, Fe i k4271, G band, Mn i þ Fe i k4458, Ca i þ Fe i k5270, Fe i k5329, Fe i k5404, and Ca i k6162. An early G spectral type for RY Tau has been found with similar (Holtzman et al. 1986) and higher resolution spectrum by Mora et al. (2001). In addition, our types for CO Ori, GW Ori, and P2441 agree within the errors with determinations by Holtzman et al. (1986) and McNamara (1990). Figure 2 shows spectra of the objects in our sample later than G2. We assign a spectral type of G6 to T Tau. This is a significantly earlier type than previous classification estimates (Herbig & Rao 1972; CK79), which many authors use. Our result is supported by the strength of features like Ca i k4226, the G band, Ca i k5270, Fe i k5404, and Ca i k6162, which are more consistent with a G5 standard than a K1 standard. T Tau shows many emission lines in its spectrum; in particular, the Mg i k5173 line conspicuous in other stars is clearly filled in. In one case, EZ Ori, we obtain differing spectral types from different parts of the spectrum. Inspection of Figure 2 indicates that the depth of the lines at wavelengths shorter than (Fe i k4047, Ca i k4226, and G band) is similar to what is observed in objects earlier than a G2 standard star, but redder features are more consistent with a G5 standard star. This could be the origin of the differing spectral types published previously, which range from F8 (CK79) to G6 (McNamara1990).WeadoptacompromisespectraltypeofG3.

6 No. 3, 2004 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1299 Fig. 2. FAST spectra of V1044 Ori, EZ Ori, T Tau, and GX Ori. Spectra of standard stars with spectral types G2, G5, and K1 are shown for comparison. The types of the targets fall within this range; see Table 1. Relevant spectral features are identified. Using our echelle spectra and the spectral types determined for the targets, we have determined the veiling at selected wavelengths 8 following the procedures of Hartigan et al. (1991). Table 5 shows these measurements and the range of wavelengths over which they were obtained. Essentially, none of the objects are veiled at within the measurement uncertainties ( ). Our results are entirely consistent with the detailed veiling study of Basri & Batalha (1990), which includes T Tau, SU Aur, RY Tau, GW Ori, and CO Ori. We also show measurements of v sin i for some of the targets in Table 5. These were determined in the usual manner by comparing our targets to appropriate standard-star templates convolved to varying levels of stellar rotation velocity. Determinations of v sin i for CO Ori and GW Ori are consistent with similar measurements in Basri & Batalha (1990; 48 and 43 km s 1, respectively). Measurements for EZ Ori and V1044 Ori have not been published before, to our knowledge. 8 We adopt the standard definition of veiling as the ratio of the excess flux to the photosphere continuum flux (Hartigan et al. 1991) Reddeningg The extinction correction at UV wavelengths appropriate to stars embedded in molecular clouds is highly uncertain. The extinction law in these regions seems to differ from that in diffuse regions in the value of total to selective reddening, R V, in the strength of the bump, and in the shape of the FUV-region extinction. Cardelli et al. (1989, hereafter TABLE 5 Veiling and Rotational Velocity Measurements Object r ( ) r ( ) v sin i (km s 1 ) T Tau SU Aur RY Tau EZ Ori V CO Ori GW Ori GX Ori

7 1300 CALVET ET AL. Vol. 128 TABLE 6 Reddening Laws k (8) HD (A k /A V ) HD (A k /A V ) Fig. 3. Comparison of extinction laws. Extinction laws for HD (Whittet et al. 2004) and HD (Clayton et al. 2003) used in this work are shown in heavy and light solid lines, respectively. For comparison, extinction laws calculated with the CCM89 formulation for R V ¼ 3:1, corresponding to interstellar reddening, and for R V ¼ 5 are also shown. CCM89) give representations of the reddening law, which depend on only one parameter, R V, that are shown in Figure 3 for R V ¼ 3:1 and 5. It can be seen that the strength of the UV extinction relative to the visual, A k =A V, decreases as R V increases, as does the strength of the bump relative to the overall extinction. In general, high-density clouds producing high values of A V tend to be characterized by high values of R V (Whittet et al. 2001, 2004). However, anomalous extinction laws have been found that cannot be represented with the CCM89 expressions. The best known example of these anomalous laws is that toward the B-type star HD 29647, located behind the Taurus clouds. This extinction law is essentially similar to the interstellar law in the optical (R V ¼ 3:6; Crutcher 1985; Whittet et al. 2001, 2004) but shows a weak bump; however, the UV extinction is considerably higher than that corresponding to high values of R V. We show in Figure 3 two of these anomalous extinction laws, the law for HD (Whittet et al. 2004) and the law for HD (Clayton et al. 2003). 9 These extinction laws are given in Table 6, and analytical fits as a function of wavelength are given in Whittet et al. (2004). Figure 4 shows spectra of a subsample of the targets corrected for reddening with three different reddening laws A k =A V : CCM89 with R V ¼ 3:1, CCM89 with R V ¼ 5, and HD To make these corrections, we need to derive the optical extinction A V from optical colors. Often, discrepant values of A V are obtained depending on whether it is derived from B V, V R C,orV I C. For low-mass TTSs, B V is strongly contaminated by excess emission, but discrepancies are also found for the other two colors. Gullbring et al. (1998) suggested that these discrepancies could be produced by extra 9 HD is a B-type star located behind the Taurus clouds, 6 0 southeast of HD emission due to stellar activity and found that V R C was the color that better represented the stellar photosphere. In the case of the IMTTSs, we found that using effective wavelengths for each filter (calculated from the V I C colors using a table kindly sent to us by M. Bessell), the discrepancy between the A V derived from different colors was greatly reduced and was within the uncertainties, 0.2, in most cases. For this tests, we compared A V derived from our colors, and also from average colors from photometry in Herbst & Shevchenko (1999), for the targets in common, using standard colors from Kenyon & Hartmann (1995, hereafter KH95). The largest discrepancy was found for the A V derived from V R C.This may be due to the fact that the ratio between the optical extinction and the color excess is largest for V R C for all

8 No. 3, 2004 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1301 Fig. 4. Effect of three extinction laws on UV fluxes of SU Aur in Taurus-Auriga (left column), EZ Ori in Ori 1c (middle column), andcooriinthek Ori ring (right column). The three laws used are interstellar law, R V ¼ 3:1(top row); R V ¼ 5(middle row); and anomalous extinction law toward HD 29647, (bottom row) (x 3.2). Overcorrection for the bump in the extinction law results in clear excess flux for the interstellar law. The effect is greatly diminished for R V ¼ 5and for the HD extinction law. extinction laws considered; for instance, for the HD law, A V =E(V R C ) ¼ 5:9 while A V =E(V I C ) ¼ 2:6. The ratio A V =E(V R C ) is also very sensitive to the change in effective wavelength with color. For these reasons, we adopted the A V derived from V I C. As shown in Figure 4, the R V ¼ 3:1 reddening correction produces a large excess around This excess is not present in stars with the lowest A V, which suggests that they are due to overcorrections by the bump at present in the extinction law. In contrast, fluxes corrected with R V ¼ 5, and especially with the HD law, are much smoother and do show much less excess. This behavior was already noticed by Herbig & Goodrich (1986) when analyzing the IUE spectra of T Tau and RY Tau; they adopted the HD law as more characteristic of the Taurus cloud. This comparison indicates that reddening laws with a weak bump are more representative of the law appropriate to these objects. However, it does not discriminate between laws with high values of R V and anomalous reddening laws. We can do so by comparing reddening corrected UV fluxes with predicted UV photospheric fluxes, fixed by the simultaneous V photometry.

9 1302 CALVET ET AL. Vol. 128 TABLE 7 IUE MK Standards Object Standard Spectral Type IUE Exposure T Tau... HD G5 LWP11114 SU Aur... HD G2 LWP12511 RY Tau... HD G2 LWP12511 EZ Ori... HD G2 LWP12511 P HD 4614 G0 LWP15018 V1044 Ori... HD G2 LWP12511 CO Ori... HD 4614 G0 LWP15018 GW Ori... HD 4614 G0 LWP15018 GX Ori... HD G8 LWP09649 The intrinsic photospheric fluxes are fixed by the V magnitude, since the veiling at V is essentially zero (x 3.1). For a given reddening law, distance, and spectral type (cf. Table 1), we calculate the stellar bolometric magnitude and luminosity L * from V, using bolometric corrections and effective temperature spectral type calibrations from KH95, from which we derive the stellar radius R *.WethenderivethemassM * interpolating between the evolutionary tracks of Siess et al. (2000), which cover the mass range of interest. We then synthesize the photospheric fluxes from the spectral library of Charlot & Bruzual (1991) using the effective temperature and gravity derived from the mass and radius. For a more detailed comparison, we use an IUE spectral standard spectrum of the same type in place of the low-resolution Bruzual-Charlot spectra in the NUV region (the STIS spectra were convolved to match the resolution of the IUE standard). Table 7 lists the IUE standard used for each target. Using this procedure, we calculate UV photospheric fluxes, fixed by the simultaneous V photometry, and compare them with the STIS NUV fluxes for two extinction laws that show a weak bump, namely, CCM89 R V ¼ 5 and HD We find that the predicted photospheric fluxes are higher than the NUV fluxes for the CCM89 R V ¼ 5 law, as shown in Figure 5 for the same stars in Figure 4. In contrast, the larger A k =A V of the HD law (see Fig. 3) results in higher NUV fluxes relative to the photospheric fluxes. Our exploration indicates that an extinction law with high values of A k =A V in the UV but a weak feature is required to explain the observations of our targets. The extinction laws of HD and HD fit these requirements. So far, a law with these properties has been considered only for the Taurus clouds (Herbig & Goodrich 1986), behind which both stars are located. However, we require similar properties of the extinction law for stars from three different star-forming regions. These stars are all located in molecular clouds, far from the destructive effects of O stars. This appears to be valid even for the stars in Ori OB1c; Figure 6 shows these stars projected against the 13 CO total intensity map of Bally et al. (1987). They are all relatively far from the highest regions of extinction, especially EZ Ori. Even the closest pair are relatively far from the cluster center (3 pcinprojected distance) where the O stars are located, which is also indicated in the figure. Fig. 5. Comparison of dereddened UV fluxes with photospheric fluxes for extinction laws that produce little overcorrection for the bump (R V ¼ 5and HD 29647; see Fig. 4 and x 3.2). Fluxes are shown for the same stars as in Fig. 4. UV fluxes are shown with heavy solid lines and photospheric fluxes with light solid lines. The photospheric flux level is fixed by simultaneous optical photometry.

10 No. 3, 2004 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1303 Fig. 6. Location of the three stars in Ori OB1c superposed on isocontours of 13 CO from Bally et al. (1987). The location of the Orion Nebula Cluster is shown by a white square. Whittet et al. (2001, 2004) argue that laws with high values of R V characterize regions with a significant lack of small grains. In contrast, laws like that of HD are representative of regions where the grain size distribution is not very different from the interstellar medium, but ice mantles have grown around grains, which tend to wash out spectral features. More recently, Whittet et al. (2004) find that the reduction of amplitude of the bump is due to an increased degree of hydrogenation of the grains brought about by the high density of the environment, while mantles tend to broaden the bump. A further condition for diminishing the bump strength is a low level of UV radiation, which would drive desorption of H atoms from the grains. As discussed, these conditions are met in the environments where our targets are located. Still, it appears that the reddening law varies in detail inside a given cloud, specially in the FUV, as shown by comparison between the reddening laws of HD and HD (cf. Figure 3). However, since these are the only two stars for which these determinations exist, we adopt their reddening laws as representative of cold molecular cloud conditions Stellar Properties As described in x 3.2, we estimate the visual extinction A V from the adopted spectral type and the E(V I C ) color excess for all targets, using the HD law. Similarly, we determine the stellar luminosity from the V magnitude and masses and ages from interpolating in the Siess et al. (2000) tracks (cf. x 3.2). These quantities are shown in Table 1. Figure 7 locates the stars in the H-R diagram, together with the Siess et al. (2000) evolutionary tracks and isochrones. We also give in Table 1 the estimated uncertainties for each quantity, which were estimated from the spectral type and color errors, and the uncertainty in extinction law. For the last, we adopted one half of the difference between magnitudes and colors dereddened with the extinction laws of HD and Fig. 7. Location of the targets on the H-R diagram, superposed on evolutionary tracks from Siess et al. (2000), labeled by their mass in solar units. Different symbols are given for each region, as indicated. (SU Aur has been shifted by 0.01 in log T ea for clarity.) Isochrones are shown as dashed lines and correspond to log age ¼ 5:5, 6 (marked ), 6.5, 7, and 7.5. The zero-age main sequence is plotted in long-dashed lines. For comparison, we plot the sample of T Tauri stars from Taurus, from KH95 (small filled circles), and the sample of Herbig Ae/Be from Hernández et al. (2004) (triangles). Thick solid lines show birth lines; from top to bottom: Palla & Stahler (1993) with mass infall rate 10 4 M yr 1, Palla & Stahler (1993) with mass infall rate 10 5 M yr 1, Hartmann (2003) with mass infall rate 10 5 M yr 1, and Hartmann (2003) with mass infall rate 2 ; 10 6 M yr 1. HD as the uncertainty. In the optical, these laws do not differ much from each other or from the interstellar reddening law, so this uncertainty does not introduce significant errors in the determination of A V and thus in L * and the other stellar properties; the errors in these quantities are mostly due to the spectral-type error Characteristics of the UV Spectra Figure 8 shows in an expanded scale the STIS/FUV and NUV spectra of EZ Ori, as representative of the sample. The spectra have been corrected for reddening using the HD law. Important absorption and emission features are labeled. The photospheric spectrum is shown for comparison. Spectra for all the targets on a smaller scale are shown in Figures The NUV spectrum is shown in the top panel. Above k , the spectrum is characterized by absorption features veiled relative to the photosphere. In particular, we mark the lines of Mg i k2852 and those of Fe ii multiplets 62 and 63, which flank the emission of the Mg ii h and k doublet. As photospheric fluxes drop toward shorter wavelengths, the excess continuum and emission lines become dominant. Lines of multiplets UV(1) to UV(3) become conspicuous in emission, and lines formed at high temperatures, hot lines (following Valenti et al. 2000, hereafter VJL00), begin to appear, as do C iii] k2334 and C ii] k2325. The FUV spectrum, shown in the lower panel of Figure 8, is dominated by the hot lines, as is the case for the lower mass CTTSs (VJL00). Table 8 lists the wavelengths of a number of

11 1304 CALVET ET AL. Vol. 128 Fig. 8. Dereddened UV spectrum of EZ Ori (solid line). The photosphere, scaled by nearly simultaneous V photometry, is shown for comparison (dotted line). Emission lines in Table 8 are identified. identified features. Many are actually a blend of several lines, which we cannot separate with our resolution. We can clearly see C ii k1776 and Si iii k1206, lines that were difficult to access in IUE spectra (VJL00). As noted by VJL00, ABWVJ02, and Herczeg et al. (2002), a large number of H 2 lines are seen in the FUV spectra of YSO, which is consistent with Ly pumping. We see them as well, although it is difficult to isolate individual lines because of the low resolution. To assess the blending, we have compared the STIS E140M spectrum of TW Hya (kindly made available by G. Herczeg) with the same spectrum convolved to the resolution of the 140L grating. In this way, we identify two strong lines as characteristic of the H 2 spectra: (1) k1504, which at the 140L resolution is still single and corresponds to line (1 7)P(5) k , part of the sequence created by cascade down from levels populated by the transition (1 2)P(5), pumped by Ly radiation at kms 1 from line center, and (2) k1524, which is a blend of lines (1 7)P(8) k , from the sequence (1 2)R(6) pumped by Lyman at km s 1, and (0 7)P(3) k , from the sequence (0 2)R(1), pumped by Ly at km s 1. These lines are marked in Figure 8. A more detailed study of the H 2 lines will be carried out in a separate paper. Table 8 gives observed line fluxes for lines identified in the STIS spectra of the targets. These fluxes were measured above the continuum, using standard IRAF routines. When a line was clearly Gaussian, we fitted it with a Gaussian function. If lines were blended, we measured the total flux between the two ends of the feature; we then fitted Gaussian curves to each line in the blend and apportioned the total flux according to these Gaussians. Measuring errors are given in Table 8; they are due to line blending and uncertainties in the continuum level adopted in the measurement. The second row in Table 8 for each line lists line luminosities, determined by correcting the observed line fluxes for reddening with the HD extinction law and using distances in Table 1. Errors listed for the luminosities combine measurement errors with uncertainties due to optical extinction and extinction-law uncertainties. We also show in Table 8 luminosities estimated from dereddened UV fluxes: L UV, with fluxes integrated between 1100 and and L FUV, with fluxes integrated between 1100 and Uncertainties have been estimated in the same way as the lines. We also list the excess UV luminosity, L exc UV,estimated by subtracting the photospheric luminosity below from the total UV luminosity. The stellar photosphere dominates in the NUV, in which most of the excess flux appears in emission lines. The only exception is the later type target GX Ori. The last row of the table gives the sum of the luminosities of the lines listed. It generally amounts to 30% 50% of the total FUV luminosity, which in turn is 30% 45% of the excess luminosity. The difference in FUV luminosity may come from H 2 emission and/or an unidentified source of continuum (x 4.1)

12 No. 3, 2004 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1305 Fig. 9. SEDs of stars in Taurus compared with best shock models with parameters in Table 9. Top: entire wavelength range, showing UV fluxes (light solid lines) and simultaneous photometry ( filled circles). The shock models are shown in heavy solid lines and the photosphere in dotted lines. Middle: NUV region. Bottom: FUV region. The FUV fluxes are the most affected by the extinction-law uncertainty, as inspection of Figure 8 shows. In fact, for stars with the largest A V, the continuum apparently rises with decreasing wavelength, cf. Figures This behavior is not expected; inspection of the same figures shows that the FUV continuum appears rather flat in stars with low reddening, as well as in the case of the low-reddening, low-mass CTTSs TW Hya and BP Tau (Herczeg et al. 2002; Bergin et al. 2003). One possible explanation is that the FUV extinction used is not appropriate. Fitzpatrick & Massa (1988) find that the curvature and strength of the FUV extinction curve show a large dispersion in dense environments and suggest that different populations of grains may be responsible for different parts of the extinction curve. Since we have no way of obtaining the best FUV extinction curve for each object, we continue with our adopted procedure but caution the reader that the uncertainties in the FUV fluxes can be larger than those we quote. 4. MAGNETOSPHERIC ACCRETION IN IMTTSs As discussed in x 1, the presently accepted paradigm for loading mass from the accretion disk onto low-mass T Tauri stars is that of magnetospheric accretion. We expect this model to be valid for stars of higher mass as well and aim to search for the specific predictions made by magnetospheric accretion. In particular, as mass from the disk flows along

13 1306 CALVET ET AL. Vol. 128 Fig. 10. Same as Fig. 9, but for Ori OB1c stars. magnetic field lines at nearly free-fall velocities, it produces emission lines with characteristic profiles (Hartmann et al. 1994; Muzerolle et al. 1998a, 1998b, 2001). In addition, as material strikes the stellar photosphere, it forms an accretion shock, which has characteristic emission especially conspicuous in the UV for late-type stars (CG98; Gullbring et al. 2000). In this section, we test these predictions against the observations. We first determine that the UV emission in IMTTSs can be interpreted as accretion shock emission. Then, we examine emission-line profiles for some of our targets and show that they are consistent with formation in the magnetospheric accretion flow. Since magnetospheric accretion can account for the bulk of the UV excess emission as well as line profiles, we are justified in using accretion shock models to calculate accretion luminosities and mass accretion rates Accretion Shock Emission The accretion shock formed on the stellar surface as freefalling material merges onto the star produces soft X-ray radiation, which heats the preshock infalling region and the photosphere immediately below the shock. The emission from the shock is essentially the emission from the optically thick, heated photosphere plus the optically thin emission from the preshock region (CG98). The emission from the shock has a higher color temperature than the stellar photosphere, thus becoming more conspicuous as wavelength decreases. For low-mass TTSs, even with small filling factors, the emission

14 No. 3, 2004 MASS ACCRETION RATES OF INTERMEDIATE-MASS TTS 1307 Fig. 11. Same as Fig. 9, but for k Ori stars. is apparent in the B and V bands and dominates in the UV (CG98; Gullbring et al. 2000; Ardila & Basri 2000). The contrast between the hotter and brighter photosphere of IMTTSs and the shock emission is lower in the optical range, making their separation difficult. The effects of the shock are expected to begin appearing in the NUV and dominate at the FUV. To seek out these effects, we have constructed accretion shock emission models with stellar parameters appropriate to each target, following the procedures of CG98. The accretion column is assumed to be plane-parallel and perpendicular to the stellar photosphere. Each column carries an energy flux F, and the columns cover a fraction f of the stellar surface. The shock emission is calculated with the Raymond (Raymond & Smith 1977) code, with the strong shock assumption. Half of this emission is used as input in the shock program CLOUDY, to calculate the preshock emission. In turn, the backward emission from the preshock region plus half of the shock emission is used as input to calculate the structure and emission of the heated photosphere. The photosphere below the shock is assumed to be in hydrostatic equilibrium, heated from above by the shock and preshock emission and from below by the intrinsic stellar flux, T 4 ea. Shock and preshock emission for a given value of the ratio M =R scale with F. This is to be expected, because the infall velocity depends on M =R. With this and stellar parameters fixed, the dependence on F isadependenceondensity(cf. CG98). So, for a given M =R foratarget,wehaveinterpolated in the shock models of CG98 and Gullbring et al. (2000)