Survey
* Your assessment is very important for improving the workof artificial intelligence, which forms the content of this project
* Your assessment is very important for improving the workof artificial intelligence, which forms the content of this project
The Astronomical Journal, 127:1682–1701, 2004 March # 2004. The American Astronomical Society. All rights reserved. Printed in U.S.A. SPECTRAL ANALYSIS AND CLASSIFICATION OF HERBIG Ae/Be STARS Jesús Hernández,1,2,3 Nuria Calvet,4,2 César Briceño,1,2,4 Lee Hartmann,4 and Perry Berlind4 Received 2003 September 16; accepted 2003 November 19 ABSTRACT We present an analysis of the optical spectra of 75 early-type emission-line stars, many of which have been classified previously as Herbig Ae/ Be (HAeBe) stars. Accurate spectral types were derived for 58 members of the sample; high continuum veiling, contamination by nonphotospheric absorption features, or a composite binary spectrum prevented accurate spectral typing for the rest. Approximately half of our sample exhibited [O i] k6300 forbidden-line emission down to our detection limit of 0.1 Å equivalent width; a third of the sample exhibited Fe ii emission (multiplet 42). A subset of 11 of the HAeBe sample showed abnormally strong Fe ii absorption; 75% of this subset are confirmed UX Ori objects. Combining our spectral typing results with photometry from the literature, we confirm previous findings of high values of total-to-selective extinction (RV 5) in our larger sample, suggesting significant grain growth in the environments of HAeBe stars. With this high value of RV , the vast majority of HAeBe stars appear younger than with the standard RV ¼ 3:1 extinction law and are more consistent with being pre–main-sequence objects. Key words: Hertzsprung-Russell diagram — stars: emission-line, Be — stars: pre–main-sequence — techniques: spectroscopic of luminosities and Teff , because it is not always straightforward to determine accurate spectral types and extinction for these objects. This results in considerable uncertainty in the location of HAeBe stars in the Hertzsprung-Russell (H-R) diagram. The presence of continua and emission lines formed outside the photosphere complicate traditional spectral classification schemes used for early-type stars. Several efforts have been made in the past to classify HAeBe stars, applying qualitative and quantitative spectral classification schemes (Strom et al. 1972; Cohen & Kuhi 1979; Finkenzeller & Mundt 1984; Finkenzeller 1985; Hillenbrand et al. 1992; Hillenbrand 1995; Mora et al. 2001). However, differences of several subclasses and even classes can be found between these various works. The discrepancies probably arise because of the different methods used. Strom et al. (1972) used solely the K line of Ca ii and He i lines to derive Teff for 18 HAeBe stars; since these lines can sometimes be found in emission, methods that rely heavily on these features cannot always be used for spectral classification. Cohen & Kuhi (1979) classified 71 H emission stars earlier than G0 using several spectral indices in the range 4270–6710 Å at a resolution of 7 Å. However, some of the features they used, such as He i kk4922, 5016 and Na i kk5890, 5896, can be affected by emission or anomalous absorption (xx 4 and 5). Finkenzeller (1985) used a scheme based on nine spectral indices in the range 3500–5000 Å, but their sample consists of only a few stars. Hillenbrand (1995) applied a quantitative spectral classification scheme to 33 HAeBe stars using features in the R and I photometric bands; however, these wavelength regions contain few useful spectral indices for classifying stars earlier than F0. Recently, in data obtained during spectroscopic campaigns carried out by the EXPORT (EXoPlanetary Observational Research Team) consortium, Mora et al. (2001) determined spectral types and luminosity classes for 29 HAeBe stars. They selected photospheric lines that did not vary on multiepoch spectra, and used high-resolution spectra to correct absorption features for rotational broadening. While this investigation yielded more 1. INTRODUCTION The Herbig Ae/ Be (HAeBe) stars are emission-line stars of spectral types B, A, and in a few cases F, in most instances spatially correlated with dark clouds or bright nebulosities (Herbig 1960; Finkenzeller & Jankovics 1984; Waters & Waelkens 1998). Comparison of their effective temperature (TeA) and luminosities with theoretical evolutionary tracks (Strom et al. 1972; Cohen & Kuhi 1979; van den Ancker et al. 1998; Palla & Stahler 1991) indicates that these objects are young, still approaching the main sequence. The HAeBe stars exhibit IR excesses, which are attributed to dust emission from circumstellar disks (Finkenzeller & Mundt 1984; Lorenzetti et al. 1983; Davies et al. 1990; Hillenbrand et al. 1992; van den Ancker et al. 1997; Malfait, Bogaert, & Walkens 1998). Millimeter observations confirm the existence of dusty disks of substantial mass around some of these objects (Mannings & Sargent 1997, 2000; Natta et al. 2000, 2001). In some cases the emission lines seen in HAeBe stars exhibit P Cygni profiles, suggesting formation in winds; more symmetric lines might arise in hot, extended chromospheres (Herbig 1960; Finkenzeller 1985; Hamann & Persson 1992). Sometimes inverse P Cygni profiles are observed, leading Sorelli, Grinin, & Natta (1996) and Muzerolle et al. (2004) to argue that the magnetospheric infall paradigm that has been applied to lowmass, accreting T Tauri stars (Muzerolle, Calvet, & Hartmann 2001) may also hold in these systems. However, despite the significant progress that has been made toward understanding HAeBe stars, problems still remain. One of these basic issues is deriving reliable estimates 1 Centro de Investigaciones de Astronomı́a (CIDA), Apartado Postal 264, Mérida 5101-A, Venezuela; [email protected], [email protected]. 2 Postgrado de Fı́sica Fundamental, Universidad de Los Andes (ULA), Mérida 5101-A, Venezuela. 3 Visiting student, Harvard-Smithsonian Center for Astrophysics. 4 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138; [email protected], [email protected], [email protected]. 1682 SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS reliable results, the extensive observational effort needed is difficult to apply to a large number of objects. Reliable spectral types are important in determining observables that can yield valuable information about the physical environment surrounding these young stars, like the value of total-to-selective extinction (RV) and infrared excesses. Some authors (Strom et al. 1972; Thé et al. 1981; Herbst et al. 1982; Sorelli, Grinin, & Natta 1996; Bibo et al. 1992; Gorti & Bhatt 1993; Waters & Waelkens 1998; Whittet et al. 2001) have suggested that HAeBe stars have higher values of RV than given by the standard interstellar extinction law (RV ¼ 3:1) frequently used to calculated the visual lineof-sight absorption (AV) toward these stars (Hillenbrand et al. 1992; Testi et al. 1998; Oudmaijer et al. 2001; Mora et al. 2001). The value of RV can be used to infer grain properties of the dust surrounding HAeBe stars. In addition, knowledge of the reddening law is essential in placing these objects in the H-R diagram and thus deriving masses and ages by comparison with evolutionary tracks. In this contribution we obtain spectral types for a large set of HAeBe stars. We use spectral indices constructed to minimize the effects of nonphotospheric emission as far as possible. In x 2 we present the observations and data reduction. The spectral classification method is described in x 3. We discuss the resulting spectral types and details of specific objects in x 4. In x 5 we discuss the anomalous features seen in our spectra. In x 6 we discuss the determination of reddening for our sample and explore the nature of the interstellar extinction law toward these objects. In x 7 we locate the stars in the H-R diagram and derive their ages and masses. A summary and conclusions are presented in x 8. 2. OBSERVATIONS Optical spectra were obtained for 75 of the 99 stars in the Herbig and Bell Catalog (HBC; Herbig & Bell 1988) having spectral types B, A, and F. Observations were made during 1999 July and 2000 January using 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 mm1 grating and a 300 wide slit. This combination offers 3400 Å of spectral coverage centered at 5500 Å, with a resolution of 6 Å. We also observed 59 main-sequence and 16 giant and subgiant standard stars covering a spectral range from O8 to M6, using the same setup (see x 3). Several spectra were obtained for most program and standard stars. The spectra were reduced at the CfA using software developed specifically for FAST COMBO observations. All individual spectra were wavelength calibrated and combined using standard IRAF routines.5 The effective exposure times for the combined spectra ranged from a few seconds to 1200 s. The signal-to-noise ratio (S/N) of our combined spectra are typically k12 at the central wavelength region of the spectra. The spectra were corrected for the relative system response using the IRAF SENSFUNC task and observations of spectrophotometric standard stars. In Figure 1 we show four examples of typical FAST spectra: one standard star, two HAeBe stars, and one ‘‘continuum’’ star (see x 4.4). 5 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. 1683 Fig. 1.—Examples of FAST Spectra of our sample. Two HAeBe stars (UX Ori and AB Aur), one continuum star (MWC 137), and a standard A0 star (HD 140775) are shown. Besides H, UX Ori does not show other emission lines; however, it exhibits anomalous absorption in the Fe ii (42) multiplet. AB Aur has H in emission, in addition to Fe ii (42), He i k5876, He i k6678, and the forbidden line [O i] k6300; there is an emission component in H. MWC 137 exhibits the entire Balmer series in emission, as well as most of the other lines. The absence of absorption features precludes the determination of a spectral type for this object. Spectra for the entire sample, plus additional information for each star, is available on-line.6 3. SPECTRAL CLASSIFICATION METHOD Spectral classification of early-type stars (B, A, and F) relies mainly on the strength of atomic absorption lines, such as the hydrogen Balmer series and He i and Fe i lines, which are sensitive to changes in TeA. In cooler objects like G stars, metallic lines start to increase in strength as a function of TeA, hence the usefulness of features such as Mg ii, Ca ii, Ca i, and the G band (CH k4300) for spectral typing. The classification scheme we present here is based on 33 spectral features that are sensitive to changes in TeA, listed in Table 1. Column (1) of this table gives an ID number for each feature band (FB), column (2) lists the main atomic/molecular species contributing to each FB, column (3) gives the central wavelength of the FB, and column (4) lists the spectral type range over which the index is useful. These spectral features were selected from previous spectral classification studies for normal stars (Morgan, Keenan, & Kellman 1943; Stock & Stock 1999; Coluzzi 1999; Gray et al. 2001; Pritchet & van den Bergh 1977; Reid et al. 1995) and studies related to HAeBe stars (Strom et al. 1972; Cohen & Kuhi 1979; Finkenzeller 1985; Waters & Waelkens 1998; Hillenbrand 1995). Following Hillenbrand (1995), the equivalent width (Wk) for each spectral feature is obtained by measuring the decrease in flux due to line absorption from the continuum that is expected when interpolating between two adjacent bands, 6 See http://cfa-www.harvard.edu/youngstars/jhernand/haebe/principal.html or the mirror site at http://www.cida.ve/~jesush/haebe/principal.html. TABLE 1 Selected Features Sensitive to Spectral Type ID (1) Features (2) kFBa (Å) (3) Spectral Range (4) Correlation CoeDcient (5) Fit Error b (6) 1................ 2................ Ca ii (K) He i He i Fe i + Sc i H H He i + Fe i He i + Fe i He i + Fe i CN + Fe i Ca i Fe i CH (G band) H H He i + Fe i He i + Fe i He i + Fe i Mn i + Fe i He i + Fe i + Mn ii He i + Fe i + Mn ii He i + Fe i + Mn ii Mn ii Fe i + Mn ii Fe i He ii + Fe i Fe i He i + Fe i He i + Fe i He i + Fe i He i + Fe i + Ti i Fe i + Ti i + Cr i Fe i + Ti i + Cr i Fe ii + Mg i Ca i + Fe i Fe i Fe i Fe i Ca i + TiO Ca i + TiO Fe i + Mg i + V i Fe i + Mg i + V i He i + Na i + TiO He i + Na i + TiO Na i + Ti i Na i + Ti i Mn i Ca i + TiO Ca i + TiO He i He i + TiO He i + TiO 3933 4026 4026 4047 4102 4102 4144 4144 4144 4175 4226 4271 4305 4349 4349 4387 4387 4387 4458 4471 4471 4471 4481 4490 4532 4669 4787 4922 4922 4922 5016 5079 5079 5173 5270 5329 5404 5404 5589 5589 5711 5711 5876 5876 5890 5890 6015 6162 6162 6678 7066 7066 A0–G0 O8–B3 B3–A0 F2–K1 O8–A1 A1–F9 O8-B3 B3–A1 F5–K3 F5–G9 F2–K3 F2–K5 F2–G2 O8–A1 A1–K6 O8–B3 B3–A1 F2–K4 F2–K4 O8–B2 B2–A1 A7–K1 B5–A1 B5–A1 A0–G5 O8–B3 A5–K3 O8–B2 B2–A1 A7–K4 A0–K5 A0–K3 K5–M5 A0–G0 A0–K0 F2–K5 O8–B2 F5–K5 A0–K3 M0–M6 A5–K5 K5–M6 O8–A0 G9–M0 F2–G2 G9–K7 F2–K5 F5–K3 K0–K7 B2–A0 O8–A1 M0–M6 0.99 0.82 0.97 0.91 0.97 0.99 0.84 0.97 0.88 0.85 0.94 0.95 0.98 0.98 0.99 0.50 0.95 0.99 0.96 0.87 –0.99 0.96 0.98 0.98 0.98 0.93 0.97 0.93 0.98 0.98 0.95 0.98 0.92 0.96 0.96 0.96 0.90 0.97 0.95 0.98 0.96 0.96 0.95 0.96 0.89 0.97 0.94 0.97 0.96 0.96 0.96 0.97 1.6 2.6 1.6 2.7 1.4 1.4 2.3 1.7 3.2 2.3 2.6 2.6 0.8 1.1 1.1 3.0 1.5 1.4 2.1 2.0 0.9 2.1 0.9 0.8 2.8 1.2 3.1 1.0 1.5 1.7 3.8 2.8 2.0 3.2 2.9 2.1 1.3 1.8 3.7 0.6 2.7 1.1 1.7 1.0 2.2 1.1 2.8 1.8 1.1 1.4 2.6 0.7 3................ 4................ 5................ 6................ 7................ 8................ 9................ 10.............. 11.............. 12.............. 13.............. 14.............. 15.............. 16.............. 17.............. 18.............. 19.............. 20.............. 21.............. 22.............. 23.............. 24.............. 25.............. 26.............. 27.............. 28.............. 29.............. 30.............. 31.............. 32.............. 33.............. a b Central wavelength of the feature band. Error obtained from our first-order fit to the index, in spectral subtypes. SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS Fig. 2.—Definition of the bands used to calculate the equivalent width. The continuum is established by interpolating at the central feature band (FB) between two adjacent continuum bands (BCB and RBC). The dashed lines indicate the boundary of each band. The dotted line shows the projected continuum. defined here as the blue continuum band [BCB] and the red continuum band [RCB]. The equivalent width Wk is defined as Wk ¼ kFB FFB 1 ; FBCB þ ðkFB kBCB =kRCB kBCB ÞðFRCB FBCB Þ ð1Þ where FFB, FBCB, and FRCB are the fluxes at the central wavelengths (kFB, kBCB, and kRCB) of the feature band and continuum bands, respectively, and kFB is the width of the FB. Figure 2 shows schematically the definition of these quantities. Equivalent widths or ‘‘indices’’ measured by this procedure are largely insensitive to reddening as long as the wavelength coverage of each band is relatively small. Indices constructed in this way should also be quite independent of the S/ N as long as the sidebands are chosen to be next to the measured feature and are wide enough to obtain a good flux estimate. Thus, judicious selection of the width of each band relies on a compromise between minimizing reddening effects and maximizing the S/ N. In order to calibrate our set of indices as a function of spectral type, we selected O8–M6 main-sequence standard stars from various lists (Garcia 1989; Gray et al. 2001; Keenan & Barnbaum 1999; Jaschek 1978; Buscombe 2001) for which we had FAST spectra. Although our sample selection was aimed at HAeBe stars, with spectral types spanning from B to F, we included later spectral types because some indices exhibit ‘‘degeneracies’’ (e.g., the G band in Fig. 3); that is, for some indices one value does not yield a unique spectral type. This approach allows us to study stars that could have large errors in their published spectral types. We measured the spectral indices in our standard star spectra and plotted them against spectral types (which were 1685 assigned a numerical scale between 18 for spectral type O8 and 75 for spectral type M5). For each index, we changed the width of the BCB, RCB, and FB (from 6 to 30 Å, in steps of 2 Å) and shifted the central wavelength of the BCB and RCB until we found the best correlation coefficient between the spectral indices and spectral types. The final value of the correlation coefficient obtained for each feature is shown in column (5) of Table 1. The width of the bands typically range from 6 to 30 Å. Once the optimum widths and central wavelengths for each band (RCB, BCB, and FB) were fixed by the best correlation coefficient, we fitted straight lines to the values of Wk as a function of spectral type within various spectral type intervals. Column (6) of Table 1 shows the error in spectral subtypes derived from these piece-wise firstorder fits for each index. Illustrative plots for four indices used in our classification scheme are shown in Figure 3. The top left panel shows the calibration for index 19 (related to He i + Fe i k4922); this index increases up to spectral type B2, where the He i k4922 line has maximum absorption; then the index decreases down to a minimum at a spectral type of A0, and increases again from F0 to roughly G9 as the Fe i k4921 line becomes strong. Index 9, corresponding to the G band, and index 23, related to Ca i k5270 (see Table 1) are shown in the top right and bottom left panels, respectively; both have a monotonic behavior. In the bottom right panel we plot the calibration for index 4, related to the H line; this index has a change in slope at spectral type A0, where the absorption in the Balmer lines is at its maximum. Our method of classification can be summarized as follows. First, we used strong, conspicuous features like the G band, Fe i, and He i lines to establish whether the star is earlier or later than A0. Then the Balmer lines H and H (Balmer indices) are used as a first guess to further narrow down the spectral type range. However, because the Balmer indices may be contaminated by emission lines and affected by luminosity effects, especially at spectral types near A0 (Morgan et al. 1943; Gray et al. 2001), these indices are not given any weight in our final determination of spectral type. Using the spectral type range guessed from the Balmer indices, we determined which other indices (not affected by effects like line/continua emission) are useful to classify the particular object. Once we have determined a specific spectral type range, we determine the spectral type for each object by computing a weighted average of the individual spectral types calculated from each index. The weights are estimated from the computed error for each index. This error has two contributions, the error from the fit to the standard main sequence as specified in column (6) of Table 1 and the error in the measured Wk. The latter is calculated by assuming Poisson statistics (Gray 1992, p. 81), such that the error in each band (FB, BCB, and RCB) is the square root of the number of counts. We then propagated the error in each band to obtain the combined measurement error in the Wk. Finally, we rejected spectral indices that yielded spectral types deviating more than 3 from the weighted average or that have an error larger than six spectral subtypes. In this way we minimized possible contamination of the indices by artifacts, emission lines, or anomalous absorption features (x 4 and x 5). This classification scheme is largely independent of luminosity because most of the indices selected are not sensitive to the surface gravity of the star. In Figure 4 we plot the spectral type determined with our method against the published spectral type for our set of main-sequence and giant standard 1686 HERNÁNDEZ ET AL. Vol. 127 Fig. 3.—Calibration for selected spectral indices. The dashed lines represent the first-order fit for each spectral index. Top left: He i + Fe i k4922; the index reaches a maximum at a spectral type B2 because of the absorption of the He i k4922 line and then diminishes up to A0 and changes slope with the the onset of Fe i k4925 absorption. Top right: Ca i + Fe i k5270; the index has a monotonic behavior from B5 to K5. Bottom left: G band; the index has a monotonic behavior in the spectral range F0–K0. Bottom right: H; the index shows a bimodal behavior. It has a positive slope from B0 to A0, following the increase in absorption in the Balmer lines, then decreases for later spectral types up to late G. stars; the overall error in our calibration does not change significantly when we include giant stars. It can be seen that uncertainties due to differences in luminosity class are small in comparison to the measurement errors. 4. SPECTRAL CLASSIFICATION 4.1. General Considerations The presence of conspicuous signs of stellar and circumstellar activity in the optical spectra of HAeBe stars distinguishes them in general from their older nonpeculiar mainsequence counterparts. One of the most characteristic features is the presence of H in emission, though some stars also exhibit an emission component in higher Balmer lines. In addition, other features can also be seen in emission among a number of HAeBe stars, such as [O i] k6300, He i kk5876, 6678, the Na i doublet located at 5890 Å and the kk4924, 5018, and 5169 lines of the multiplet 42 of Fe ii (e.g., MWC 137 in Fig. 1). Some HAeBe stars show absorption features in their spectra that appear anomalous when compared with main-sequence stars of the same spectral type. Examples are the Na i kk5890, 5896 doublet and the multiplet 42 of Fe ii in absorption (e.g., UX Ori in Fig. 1). These anomalous absorption features are believed to be caused by material surrounding the star. Because the emission as well as the anomalous absorption lines (x 5) are thought to originate outside the stellar photosphere, we collectively call them nonphotospheric features. Continuum radiation generated outside the stellar photosphere can affect the stellar flux in the bands used to determine a spectral type (Herbig 1960; Hamann & Persson 1992; Corcoran & Ray 1997; Böhm & Catala 1994). The superposition of a nonphotospheric continuum on the stellar spectrum (veiling) is an effect that has to be taken into account when attempting to classify pre–main-sequence (PMS) stars. Veiling reduces the depth of absorption features and is wavelength dependent, affecting the spectral type determination. Some of the stars in the HAeBe sample show almost all features in No. 3, 2004 SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS 1687 Out of the 75 stars for which we have spectra, we have determined spectral types for 58 objects, or 77% of the sample. Of these 58 stars, 46 have a spectral type earlier than F7. Within this early spectral type sample, 39 are cataloged as HAeBe (see x 4.2); the remainder have an uncertain evolutionary status and are discussed in x 4.3. There are 12 stars that we classified as F7 or later and discuss in x 4.4. The spectra of seven stars were too veiled for spectral classification; these are discussed in x 4.5. Finally, we could find no consistent spectral type from the various indices for 10 objects of the sample, which are discussed in x 4.6. 4.2. Herbig Ae/Be Stars Fig. 4.—Comparison of spectral types determined in this work with those from the literature for main-sequence (dots) and giant (triangles) standard stars. The points follow well a straight line of slope unity, with small scatter, indicating that our calibration is largely independent of surface gravity (for luminosity classes V and III.) Error bars indicate the uncertainty obtained from our classification scheme. emission (Herbig 1960; Hamann & Persson 1992; Waters & Waelkens 1998). These objects (sometimes called continuum stars because of the presence of strong, nonphotospheric continuum emission) cannot be classified because of the almost complete absence of photospheric absorption features (e.g., MWC 137 in Fig. 1). The spectral classification scheme described in x 3 is designed to largely avoid problems caused by nonphotospheric (emission and absorption line) contributions. We achieve this by relying on many indices that are sensitive to TeA and requiring that the various spectral types calculated from each index agree with the others; wildly discrepant values are rejected, and a weighted mean spectral type is obtained. In this way, anomalous values for indices such as Ca ii k3933 (index 1), He i + Fe i k4922 (index 19), He i + Fe i + Ti i k5016 (index 20), and Na i + Ti i k5890 (index 29), which could be affected by a nonphotospheric contribution, can be detected and not included in the calculation of the final spectral type. Although stellar rotation can in principle affect line indices over narrow bandpasses (Mora et al. 2001; Gray et al. 2001), it is unlikely to bias our determinations because our bandpasses are wide; our spectral resolution 300 km s1 at 6000 Å is larger than the typical rotational velocities of HAeBe stars, 225 km s1 (Finkenzeller 1985; Böhm & Catala 1995). A further complication can arise when a star is a member of a spectroscopic binary system. The spectroscopic binary frequency for HAeBe stars is larger than 35% (Corporon & Lagrange 1999), so there is a finite possibility of observing a combined spectrum in our sample. However, the primary component tends to be more luminous than its companion (Corporon & Lagrange 1999), so we expect that the indices we use to classify the stars are probably dominated by features of the primary star in most cases. Still, there are a few likely cases of composite spectra. In Table 2 we show the spectral types determined for the 39 HAeBe stars in our sample. About 80% of the stars have an error of P2.5 spectral subtypes. Stars with high reddening tend to have larger errors because high extinction makes it difficult to measure indices at short wavelengths. In Figure 5 we compare spectral types derived in this work with those published previously for different subsets of our entire sample (Cohen & Kuhi 1979; Finkenzeller 1985; Hillenbrand 1995; Mora et al. 2001). Our determinations correlate best with those from high-resolution spectra by Mora et al. (2001; x 1). Only the HAeBe star VV Ser (HBC 282) differs by more than three subclasses in this comparison. The presence of lines of He i kk4026, 4144, 4387, 4471, and 5876 favors our determination. Mora et al. (2001) cite a large uncertainty in their determination (more than five subclasses), perhaps in part due to difficulties in typing at high resolution for a rapidly rotating star (Vrot > 229 km s 1). There are large discrepancies in the spectral types determined for some stars. The spectral types published for LkH 208 (HBC 193) range from B5 to F0; our result based on the Ca i kk5270, 5589, Fe i k5079, and Mg i k5711 lines is A7 3 subclasses. We find in our spectra that some indices are contaminated by nonphotospheric Fe ii (42)7 absorption features. This could explain the large discrepancies between various authors. The star LkH 234 (HBC 309), cataloged by Cohen & Kuhi (1979) as O9, shows enough neutral helium lines in its spectrum to favor a spectral type B7 3.5 subclasses, which is in better agreement with Mora et al. (2001) and Finkenzeller (1985). Stars LkH 338 (HBC 196) and LkH 339 (HBC 197) show the largest discrepancies between the spectral types given by Cohen & Kuhi (1979) and our determination. Cohen & Kuhi (1979) classified LkH 338 and LkH 339 as F2, but our spectra show no evidence of the G band or metallic lines expected if this were the appropriate spectral type. 4.3. Objects with Uncertain Evolutionary Status In this subsection we analyze objects that need further study in order to clarify their evolutionary status. In Table 3 and Figure 5 we compare spectral types derived in this work with those published previously for different subsets of our entire sample (Cohen & Kuhi 1979; Finkenzeller 1985; Hillenbrand 1995; Mora et al. 2001). Some stars have already been rejected by some authors as members of the HAeBe class. LkH 341 (HBC 201) was rejected by Thé et al. (1994) because of the absence of excess at far-infrared (FIR) bands. This star was classified by Cohen & Kuhi (1979) as B3, but in our spectra we clearly detect metallic lines including 7 The multiplet number of the element is given in parentheses. TABLE 2 Stars Classified as HAeBe HBC 3.......... 78........ 154...... 170...... 192...... 193...... 196...... 197...... 219...... 282...... 284...... 293...... 305...... 309...... 310...... 313...... 324...... 329...... 334...... 348...... 350...... 373...... 430...... 451...... 464...... 492...... 493...... 528...... 529...... 548...... 551...... 686...... 689...... 705...... 726...... 730...... 734...... 735...... 736...... Name V633 Cas AB Aur T Ori RR Tau HD 250550 LkH 208 LkH 338 LkH 339 V590 Mon VV Ser AS310 NW PX Vul LkH 324 LkH 234 BD +46 3471 LkH 233 MC 1 VX Cas RNO 6 IP Per XY Per EW V892 Tau UX Ori HD 245185 CQ Tau p26887 V350 Ori LkH 215 HD 259431 LkH 218 LkH 220 WW Vul V1685 Cyg LkH 147 HD 200775 BD +65 1637 BH Cep BO Cep SV Cep Spectral Type B9 A1 A0 A0 B9 A7 B9 A1 B7 B6 B1 F3 B8 B7 A0 A4 A7 A0 B3 A6 A5 B8 A3 A1 F3 A6 A1 B6 B6 A0 B8 A3 B3 B2 B3 B4 F5 F4 A0 Error 2.5 1.5 2.5 2.0 1.5 3.0 3.5 3.0 2.0 2.0 2.0 1.5 2.5 3.5 1.0 3.0 2.5 1.5 2.5 2.0 1.5 3.0 2.5 2.0 2.0 3.0 2.5 2.5 2.5 2.0 2.0 2.0 2.0 3.5 1.0 1.0 2.0 1.0 1.5 Str72 ... B9 ... A3 ... B8 ... ... A2 ... ... ... ... B5 A2 A7 a ... ... ... ... ... ... ... ... ... ... B7 B5 ... ... ... B2 ... B5 B2 ... ... ... CK79 Fink84 Fink85 Hill95 Mora01 B3 A0 A5 A6 B6 F0 F2 F2 B9 a Ae B9 A3 A3–A5 B9 B5–B9 ... ... B8 B1–B3 B–A ... ... B5–B7 A0 A7 ... ... ... ... A2II–B6 A0 ... A5 ... ... ... B7–B8 B6 B6 B5 ... B2 ... B3 B5 ... ... ... A5 A0 A2 A3 B9 A2 ... ... B7 B9 ... ... ... B3 A0 A5 ... ... ... ... ... A6 A3 A1 ... ... ... B7 B5 A0 ... ... B2 ... B3 B3 ... ... ... ... A0 A5 a A0 a ... ... a a a B0 F5 B5 O9 ... A7 ... ... ... ... ... ... ... ... ... ... ... B1 A0 B6 B5 ... ... a a ... ... ... ... ... ... ... a A4 a A5 ... ... ... a a ... A0 ... ... ... a B2 a a ... B3 ... B2.5 B3 ... ... ... a A3 IV A0 IV ... ... ... ... ... A0 V ... F3 V ... B5 V ... ... ... A0 V ... ... A2 IV ... A4 IV ... F5 IV ... A2 IV ... ... ... ... A2 IV B2 ... ... ... F5 III F5 V A2 IV Notes.—(Str72) spectral type from Strom et al. 1972; (CK79) spectral type from Cohen & Kuhi 1979; (Fink84) spectral type from Finkenzeller & Mundt 1984; (Fink85) spectral type from Finkenzeller 1985; (Hill95) spectral type from Hillenbrand 1995; (Mora01) spectral type from Mora et al. 2001. a Star observed but spectral type not assigned. TABLE 3 Objects with Uncertain Evolutionary Status HBC Name Spectral Type Error Str72 CK79 Fink84 Fink85 Hill95 Mora01 7.......... 160...... 201...... 281...... 297...... 314...... 482...... LkH 201 PQ Ori LkH 341 LkH 118 V751 Cyg LkH 350 BN Ori B2 F3 F3 B1 A0 B8 F4 2.5 1.5 2.5 2.0 2.5 3.0 2.0 ... ... ... ... ... B5 ... B3 F5 B3 ... ... ... ... ... ... B5 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... a a A5 ... ... ... ... ... Notes.—(Str72) spectral type from Strom et al. 1972; (CK79) spectral type from Cohen & Kuhi 1979; (Fink84) spectral type from Finkenzeller & Mundt 1984; (Fink85) spectral type from Finkenzeller 1985; (Hill95) spectral type from Hillenbrand 1995; (Mora01) spectral type from Mora et al. 2001. a Star observed but spectral type not assigned. SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS 1689 Fig. 5.—Comparison of spectral types determined in this work for HAeBe with previously published values. References are shown in Tables 2 and 3. Vertical error bars are the uncertainties derived from our spectral-type determination as explained in the text. For comparison, we show in each panel the line with slope 1. The largest scatter is observed when comparing our results with Cohen & Kuhi (1979). The best correspondence is obtained when comparing our spectral types with those of Mora et al. (2001). Ca i kk4226, 5270, 5589, Fe i kk4387, 4922, and Mg i k5711, in addition to the G band, indicating a later spectral type. Our analysis from all these indices yield the same spectral type F3 2.5 subclasses. Using near-IR photometry from Cohen & Kuhi (1979), we derive for LkH 341 a small near-IR excess using a standard extinction law (RV ¼ 3:1), but this excess disappears if we instead use an extinction law with RV ¼ 5:0 (see x 6). This behavior is also observed in the star LkH 118 (HBC 281) when using JHKL magnitudes from de Winter et al. (2001). This object was also rejected by Thé et al. (1994) as HAeBe. The spectral type for LkH 118 differs by more than four subclasses from that given by Finkenzeller (1985). We based our result on the He i and He ii lines. LkH 201 (HBC 7) and LkH 350 (HBC 314) were cataloged by Herbig & Bell (1988) as possible background Be stars and rejected by Thé et al. (1994) as HAeBe because of the absence of excess at FIR bands. However, these stars show some characteristics typical of PMS stars. We found emission at H , H, and Fe ii (37, 38, 40, 42, 49, and 74) in LkH 201 (x 5). Similarly, LkH 350 exhibits H and H in emission, in addition to some abnormal absorption features due to the diffuse interstellar bands (DIB), located at 5780, 5796, and 6283 Å, and to the Na i doublet (k5890; Miroshnichenko et al. 2001). However, emission in the Balmer lines and Fe ii can also be found in more evolved Be stars (Miroshnichenko et al. 2003). Both stars exhibit high reddening (AV > 5) and an anomalous extinction law (RV > 3:1, x 6), which would be expected if these stars are embedded in a molecular cloud. It is not clear if BN Ori (HBC 482) is a PMS object. It was rejected by Thé et al. (1994) as an HAeBe object because of the lack of FIR excess. We do not detect excess at J, H, and K bands, using Two Micron All Sky Survey (2MASS) photometry. However, some studies of this star propose a PMS status, suggesting that it could be a UX Ori object (Marconi et al. 2001) or that it has experienced an FU Ori type outburst (Shevchenko et al. 1997). PQ Ori (HBC 160) does not show significant differences from an F3 main-sequence star. In addition, it lacks emission lines and near-IR excess, as determined from 2MASS photometry. V751 Cyg (HBC 297) has been cataloged as a 1690 HERNÁNDEZ ET AL. cataclysmic variable by Robinson (1973), Downes et al. (1995), and Echevarrı́a et al. (2002). This star falls below the main sequence in the H-R diagram (x 7). One characteristic of these later-type objects is that the Li i k6708 absorption line seems to be present in most of these stars. Li i has been used in the past as an indicator of youth in intermediate- and late-type stars (Strom et al. 1989). However, we caution that Li i k6708 in absorption cannot be taken as an indicator of the PMS nature of stars earlier than mid-K, because the shallow depth of the convective zone in these stars can allow them to reach the main sequence with a nonnegligible amount of their primordial lithium content (Briceno et al. 1997). Therefore, the presence of lithium in absorption (Wk > 0:1 Å) in this spectral type range is only evidence that these objects are not old disk stars. 4.4. Stars with Spectral Types Later than F In Table 4 we list the 12 stars of our sample with spectral types F7 or later. Although they appear in the HBC (Herbig & Bell 1988) as earlier than G0, our classification scheme yielded types as late as G4 for some of them. For V1686 Cyg there is no agreement on the spectral type assigned by different authors, the values ranging from B2 (Thé et al. 1994) to F2 (Terranegra et al. 1994). In particular, Mora et al. (2001) assigned a spectral type A4 with more than five subclasses of spectral type error. In their multiepoch spectra, kindly provided to us by B. Merin, absorption features tend to vary significantly in time, which could explain why a reliable spectral type is rather difficult to determine. The strongest DIBs (kk 5780, 5797, 6284, and 6614) are clearly seen in the EXPORT spectra. However, our spectra look very different, the presence of the G band and metallic lines (Fe i, Ca i, Ca ii, and Mn i) are more consistent with a spectral type F9. This star exhibits large photometric variations. The brightness of V1686 Cyg decreased progressively by more than 4.5 mag in a period of 7.5 yr, then it brightened by 4 mag in about 4 yr. The decrease in brightness was accompanied by a reddening of the star (in VI ), suggesting that it could be caused by dusty material not too far from the star. In addition to this longterm variation, V1686 Cyg shows changes in brightness of more than 2 mag on timescales of roughly 2 months. These shorter term photometric variations could be related to the spectroscopic variability, but this remains to be investigated; this object deserves further study to clarify the physics involved in its behavior. The spectra of the stars LkH 349 (HBC 308) and RNO 63 (HBC 518) show P Cygni profiles at H , indicating ejection of material at velocities larger than 300 km s 1 (see also Hessman et al. 1995). In VSB 2 (HBC 531) H is seen in absorption, but its Wk is smaller than in a standard star of the same spectral type, which may be the result of unresolved emission. Among our sample of later type objects, only W84 (HBC 217) and V360 Mon (HBC 231) show Wk of H in emission k10 Å and emission in H. The [O i] k6300 and Fe ii (42) lines were not detected in any of the objects listed in Table 4. 4.5. Continuum Stars In a subset of objects we found essentially no absorption features at our resolution, so they could not be assigned a spectral type. Most of the lines appear in emission. These stars are the continuum stars, and they are listed in Table 5. In Table 6 we present measurements of the Wk of emission lines seen in these stars. Previous attempts to assign spectral types to these stars are given in Table 5, but they should be treated with caution given their high degree of veiling. The stars MWC 1080 and PV Cep show strong P Cygni profiles in several Balmer lines that are resolved even at our low resolution; this suggests the presence of strong winds or outflows. High-resolution spectroscopy of some of these stars (Fernandez 1995; Corcoran & Ray 1997; Parsamian et al. 1996; Magakian & Movsesian 2001) confirm the P Cygni nature of the line profiles. In addition, they tend to be associated with strong molecular outflows and/or optical jets (Wu et al. 1996; Arce & Goodman 2002; Magakian & Movsesian 2001; Gomez et al. 1997) pointing to the youth of these objects. Recent high-resolution observations of LkH 101 show extremely peculiar double line profiles, unlike those found in any other HAeBe star; this suggests that LkH 101 may not belong to the HAeBe class (G. Herbig 2003, private communication.) 4.6. Stars with Unknown or Uncertain Spectral Types We found cases in which it was impossible to assign a unique spectral type to the object. These objects are discussed below individually and are listed in Table 7, together with previous spectral type determinations that have appeared in TABLE 4 Stars with Spectral Types Later than F7 HBC Name Spectral Type Error Str72 CK79 Fink84 Fink85 Hill95 Mora01 217............ 222............ 231............ 308............ 432............ 436............ 442............ 460............ 518............ 531............ 535............ 690............ W84 W108 V360 Mon LkH 349 P102 RY Ori P1394 MV Ori RNO 63 VSB 2 W121 V1686 Cyg F7 F7 G4 F9 F7 F7 F8 G1 F7 G1 G2 F9 2.0 2.0 2.5 2.0 2.0 2.5 1.5 3.0 2.5 2.0 2.0 3.5 ... ... ... ... ... ... ... ... ... ... ... ... F8 F9 F8 F8 ... ... ... ... F6 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... B5 ... ... ... ... ... F5 IV ... ... ... ... ... A4 V Notes.—(Str72) spectral type from Strom et al. 1972; (CK79) spectral type from Cohen & Kuhi 1979; (Fink84) spectral type from Finkenzeller & Mundt 1984; (Fink85) spectral type from Finkenzeller 1985; (Hill95) spectral type from Hillenbrand 1995; (Mora01) spectral type from Mora et al. 2001. TABLE 5 Continuum Stars HBC Name Str72 CK79 Fink84 Fink85 Hill95 Mora01 40........... 164......... 199......... 207......... 317......... 330......... 696......... LkH 101 V380 Ori MWC 137 R Mon MWC 1080 V594 Cas PV Cep ... B8 ... ... ... B8 ... C B9(C) ... B0 B0(C) O9.5 ... ... A1e Cont + e e+s eq B8, B9eq ... ... ... Cont + e ... ... B8eq ... ... B9 B0 B0 B0 B8 ... ... ... ... B8 IIIev ... ... ... Notes.—(Str72) spectral type from Strom et al. 1972; (CK79) spectral type from Cohen & Kuhi 1979; (Fink84) spectral type from Finkenzeller & Mundt 1984; (Fink85) spectral type from Finkenzeller 1985; (Hill95) spectral type from Hillenbrand 1995; (Mora01) spectral type from Mora et al. 2001. TABLE 6 Emission Lines in Continuum Stars Element Ca ii (1) .................................. H + Ca ii (1) ......................... He i (18)................................. H........................................... Fe ii (27, 28) .......................... Fe ii (27) ................................ Fe ii (27) + Ti ii (41) ............. Ti ii (41) ................................. H .......................................... Fe ii (27) ................................ Fe ii (27) + He i (51) ............. Ti ii (19) ................................. Fe ii (27) ................................ Ti ii (19) ................................. Ti ii (31) ................................. He i (14)................................. Fe ii (37) ................................ Fe ii (38) ................................ Fe ii (37, 38) .......................... Fe ii (37) ................................ Fe ii (38) ................................ Fe ii (37) ................................ Fe ii (38) ................................ Fe ii (37, 38) .......................... Fe ii (38) ................................ Fe ii (37) ................................ Fe i (37) ................................. Fe i (43) ................................. Cr ii (30) ................................ H .......................................... Fe ii (42) ................................ Fe ii (42) ................................ Fe ii (42) ................................ Ti ii (70) ................................. Fe ii (49) ................................ Fe ii (49) ................................ Fe ii (49) ................................ Fe ii (48) ................................ Fe ii (49) + Cr ii (43)............. Fe ii (48, 49) + Cr ii (43) ...... Fe ii (49) ................................ Fe ii (48) ................................ Fe ii (48) ................................ Fe ii (49) ................................ Fe ii (55) ................................ [N ii] (3) ................................. k (Å) 3934 3969 4026 4102 4176 4233 4301 4313 4340 4352 4385 4395 4417 4445 4468 4472 4491 4508 4521 4534 4549 4556 4576 4584 4621 4629 4667 4731 4824 4861 4924 5018 5169 5189 5198 5235 5255 5264 5276 5317 5326 5338 5363 5425 5535 5755 LkH 101 ... ... ... 6.7 ... ... ... ... 9.0 ... ... ... ... ... ... ... ... ... 2.5 ... 2.3 2.6 ... 3.7 ... 2.3 ... ... ... 37.7 3.1 4.3 1.8 ... 2.2 1.4 ... ... 3.2 4.7 ... ... ... ... 1.6 ... V380 Ori 6.5 4.6 ... 0.7 4.8 2.7 6.9 ... 3.6 ... 1.29 1.5 2.0 1.4 1.1 ... 1.6 ... 5.0 ... ... 5.8 ... 4.9 1.9 3.7 1.1 0.9 0.9 12.1 6.8 8.0 9.1 1.7 3.0 4.1 ... ... 5.3 6.7 0.6 1.0 2.3 1.7 2.6 ... MWC 137 ... 5.6 1.4 7.9 1.6 1.1 ... ... 14.9 ... 0.9 ... ... ... ... 1.8 1.0 ... 1.3 ... ... 2.1 ... 1.7 ... 1.5 ... ... ... 54.4 2.4 3.4 1.7 ... 0.9 1.0 ... ... 1.6 2.4 ... ... 0.8 ... 0.7 0.7 R Mon 1.8 ... ... ... 2.1 2.0 3.7 ... 1.7a 2.0 0.6 0.8 1.3 ... ... ... 1.0 ... 2.4 ... ... 2.7 ... 2.5 0.5 1.6 0.3 0.6 ... 12.7 3.5 3.9 4.5 1.0 1.7 2.1 ... 1.1 2.9 3.1 ... 0.4 1.0 0.5 1.4 ... MWC 1080 a 0.7 0.9a ... 1.9a 3.0 2.0 3.0 0.8 4.5a 1.9 1.3 0.9 0.9 ... 0.6 ... 1.3 0.7 3.0 0.4 1.8 2.9 0.7 3.3 0.9 2.2 0.4 0.4 0.5 20.9a 4.0a 4.9a 5.0a 0.8 2.0 2.5 0.4 0.5 3.0 4.2 ... 0.5 1.1 0.6 1.1 ... V594 Cas PV Cep 0.9 1.5 ... 1.5a 1.0 1.0 1.8 ... 2.9a ... 0.7 0.7 ... ... ... ... 0.3 0.4 1.2 ... 1.3 1.0 0.2 1.6 0.2 0.8 ... ... ... 7.5a 2.8 3.8 4.2 0.4 0.7 1.28 ... ... 1.4 2.1 ... ... 0.6 0.3 0.6 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 15.2 2.5 3.0 5.0 ... 3.2 ... ... 2.0 ... ... 3.4 ... ... ... ... ... 1692 HERNÁNDEZ ET AL. Vol. 127 TABLE 6—Continued Element k (Å) Ca i (47)? ............................... He i (11)................................. Na i (1)................................... Cr i (7) + Si ii (4)?................. Cr i (7)?.................................. Fe ii (74) ................................ Fe ii (74) ................................ Fe ii (74) + Si ii (2)................ [O i] (1).................................. Fe i (1016)?............................ Si ii (2) ................................... [O i] (1).................................. Fe ii (40) + Si ii (2)................ [V i] (13)? .............................. Fe ii (74) ................................ Fe ii (40) + [V ii] (13) ? ........ Fe ii (74) + [V ii] (13) ? ........ Ca i (18)? ............................... Fe ii (40) ................................ [N ii] (1) ................................. H .......................................... He i (46)................................. [S ii] (2).................................. [S ii] (2).................................. He i (10)................................. 5857 5876 5892 5979 5992 6149 6238 6248 6300 6317 6347 6363 6370 6382 6417 6433 6456 6492 6516 6548 6563 6678 6717 6731 7066 a LkH 101 1.3 3.8 ... 1.6 ... ... ... 2.4 3.9 4.4 2.8 1.2 1.6 4.1 ... 0.3 2.2 2.2 1.2 ... 464.1 1.5 ... ... 2.8 V380 Ori 0.4 0.6 2.9 ... 0.8 2.2 2.1 3.0 ... 0.6 3.0 ... 1.3 0.4 1.0 1.3 3.64 0.7 2.7 ... 75.3 0.5 ... ... ... MWC 137 ... 5.6 ... 0.3 ... ... ... 1.2 0.8 1.1 0.8 ... 0.6 ... 0.2 0.3 0.6 0.6 0.9 ... 397.0 2.2 ... ... 2.7 R Mon ... ... 0.6 ... 0.2 0.8 ... 1.3 3.4 0.4 ... 1.6 ... 0.5 0.4 0.4 1.1 ... 0.6 ... 106.8 ... 0.6 1.1 ... MWC 1080 ... ... ... ... 0.3 1.0 0.8 1.3 0.4 0.7 ... ... ... 1.0 0.5 0.6 1.6 ... 1.0 6.2 135.9 ... ... ... ... V594 Cas ... ... 0.6 ... ... 0.4 0.3 0.6 0.2 0.3 0.4 ... 0.3 0.3 0.2 0.3 0.8 ... 0.3 ... 67.9a ... ... ... ... PV Cep ... 0.8 ... ... ... ... ... 2.6 66.4 ... ... 21.6 ... ... 1.4 2.2 1.5 1.5 2.1 4.5 125.0 1.0 7.7 14.4 1.0 P Cygni profile. the literature. Short comments about each star as well as a rough estimate of spectral type, based on any absorption features visible, are given in the last column of the table. MacC H12.—No reliable spectral type could be obtained for this star. The high reddening seen toward this star precludes the use of the blue region of the spectrum for spectral typing, especially at wavelengths below 5500 Å. However, the G band and Fe i k5329 seem to be present, which suggests a spectral type around F4. HK Ori.—This star has multiple emission lines, including Fe ii (27, 37, 67, 42, 48, and 49), Ti ii (41 and 69), Cr i (31), and Cr ii (43), which contaminate most of the indices located at wavelengths below 5500 Å. In contrast, between 5500 and 6300 Å the spectrum seems to be free of emission features. The spectral type we derive from this wavelength region is F2, based on the Ca i k5589, Fe i + Mg i k5711, and Mn i k6015 indices. However, Ca i k6162 is more consistent with a G0 star, while the Ca ii K line is similar to that expected in an A2 main-sequence star. This behavior was already reported by Strom (1983), who found a variation of spectral type ranging from early A near 4000 Å to late F around 6500 Å. One possibility is that we are observing a combined spectrum (A star + F star); the multiplicity of this star is well known (Leinert et al. 1997; Pirzkal et al. 1997; Corporon & Lagrange 1999), and the presence of Li i k6707 suggests that the companion is likely to be a star with a spectral type later than F7 (Wk [Li i] = 0.2 Å) BF Ori.—When we attempt to derive a spectral type for this star using the indices in Table 1, we obtain an unreasonably large error, given the quality of the spectrum. The indices TABLE 7 Stars Not Classified HBC Name Str72 CK79 Fink84 Fink85 Hill95 Mora01 Comments 1................. 94............... 169............. 202............. 273............. 321............. 325............. 716............. 717............. 742............. MacC H12 HK Ori BF Ori VY Mon KK Oph MacC H4 V376 Cas V1493 Cyg LkH 168 MacC H1 ... B7 ... ... ... ... ... ... ... ... A5–F A4 A0 O9 ... A9 B5 A2 F2 B8:e ... A–F ... ... A5–A7 ... ... ... ... ... ... A4 A–F ... A5–A7 ... ... ... ... ... ... A3 A7 ... A3 ... F0 ... ... ... ... G1 V A2 IV A5 Vep A8 V ... ... ... ... ... High reddening; F4? Emission lines; binary system; A2–G0 Nonphotospheric absorption lines; A0A5 Emission lines + High reddening ; F2 ? Emission lines; binary system; A0–F0 High reddening; earlier than A5 High reddening; A3–F2 High reddening; Fe ii (42) in absorption; A1–A9 High reddening; F0 ? High reddening + emission lines; B2A0 Notes.—(Str72) spectral type from Strom et al. 1972; (CK79) spectral type from Cohen & Kuhi 1979; (Fink84) spectral type from Finkenzeller & Mundt 1984; (Fink85) spectral type from Finkenzeller 1985; (Hill95) spectral type from Hillenbrand 1995; (Mora01) spectral type from Mora et al. 2001. No. 3, 2004 SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS Fe i + Ti i k5079, Ca i + Fe i k5270, and Ca i k5589 indicate that BF Ori has a spectral type A0–A9. However, when compared with an A7 standard, this star shows strong abnormal absorption in some features, which contaminate most of the spectral indices (x 5). VY Mon.—The large reddening (AV > 7:0; Casey & Harper 1990; ), together with the presence of emission lines, results in a large uncertainty in the spectral type we obtain from indices at wavelengths less than 5500 Å. The indices Ca i k5589, Mg i k5711, and Mn i k6015 yield a spectral type F2 5 subclasses, which is consistent with the presence of the G band. KK Oph.—The emission lines present in this object are so numerous that they affect most of the indices we have defined in Table 1 (x 5). Binarity is reported in this star by Bailey (1998), Leinert et al. (1997), and Pirzkal et al. (1997). The companion is probably a T Tauri star; however, in contrast to HK Ori, no Li i absorption is seen in our spectrum in spite of a good S/ N. The absence of the G band, and of the Ca i k6162, Ca i k5589, Mn i k6015, Mg i k5711, He i k6678, and He i k7066 lines suggest a spectral type between A0 and F0. MacC H4.—This star has a large extinction, which results in a poor S/N at the blue end of our spectrum, precluding an accurate determination of spectral type. Some helium lines are observed marginally at wavelengths less than 5500 Å. The Ca i, Fe i, and Mg i are clearly absent in the red part of the spectrum; this could indicate a spectral type earlier than A5. Emission components are observed in the lines He i kk5876, 6678, [O i] k6300, and the Balmer lines, H, H, and H (x 5). V376 Cas.—Indices Fe i k4532, He i + Fe i k4922, Mg i k5173, Ca i k5270, Ca i k5589, and Fe i + Mg i k5711 lead to a spectral type A8, with a large uncertainty of five subtypes. The Wk of H and H are characteristic of stars with spectral types around A5. Indices with wavelengths below 4500 Å could not be used because of the large reddening in our spectrum. This object has the largest linear polarization observed so far in any PMS star. Although some authors have suggested that the reflected light can be produced by a circumstellar disk observed nearly edge-on (Asselin et al. 1996; Hajjar & Bastien 2000), the morphology of the reflection nebula is completely different from other edge-on disk systems, resembling instead that of objects with outflow holes, as discussed by Whitney & Hartmann (1993). Indeed, Hajjar & Bastien (2000) argue that this object is an extreme Class I object, i.e., a protostar with an opaque infalling envelope. Because the star is not observed directly, but only in scattered light, reddening corrections and thus luminosity estimates are extremely uncertain. This object and the HAeBe star V633 Cas are the brightest objects in the isolated molecular cloud L1265. V1493 Cyg.—The high reddening of this star and the presence of nonphotospheric absorption features (x 5) complicate attempts to determine a reliable spectral type. Still, a spectral type A1–A9 is derived from the indices Ca ii k3933, Fe i k4787, Fe i k5079, Ca i + Fe i k5270, and Ca i k5589. Weak emission is detected in the forbidden line [O i] k6300. LkH 168.—Published spectral types range from A3 (Fernandez et al. 1995) to F6 (Terranegra et al. 1994). Because of the high reddening, the blue part of the spectrum is noisy, but the absence of the G band and the weakness of the metallic lines in the red part of the spectra indicate that the spectral type is probably earlier than F0. According to Herbig & Bell (1988) this object may be a background Be star. MacC H1.—Lines He i kk4387, 4471, 5876, and 7066 indicate a spectral type between B2 and A0, and the absence of 1693 the G band seems consistent with this estimate, but the high reddening and the presence of numerous emission lines do not allow us to obtain a more reliable spectral type. This object was included by Thé et al. (1994) as an emission-line star but not considered as HAeBe. However, we observe emission in the Balmer and Fe ii (38, 37, 42, and 49) lines in addition to P Cygni profiles in H and H. More data are necessary to study the evolutionary status of this object. 5. NONPHOTOSPHERIC FEATURES As already mentioned, HAeBe stars exhibit a number of spectral features in emission or absorption, not seen in standard stars of the same spectral type, that suggest their origin is outside the stellar photosphere. In Table 8 we list these nonphotospheric features measured in our sample, together with their Wk. In this table we include a footnote indicating the form of the H profile in our spectra. When asymmetries are seen, they could be produced by material moving at velocities larger than our spectral resolution (300 km s1). However, given our low spectral resolution, we cannot say anything conclusive about the shape of the lines. All stars, by definition of the class, show H in emission. The distribution of the Wk (H ) is shown in the top panel of Figure 6, where HAeBe stars, continuum stars, and stars with spectral types later than F are shown separately. It can be seen that the continuum stars have the largest Wk (H) suggesting that they are the youngest of the sample, since activity, powered either by disk accretion or by stellar dynamos, is expected to decrease with age (Hartmann et al. 1998; Skumanich 1972). Among the sample, 53% show emission in H and 15% in H, although these are lower limits, since at our resolution we could not detect an emission component superimposed on an absorption profile. However, we find that among the 39 HAeBe stars in Table 2, 95% have Wk (H) and 56% have Wk (H) smaller than that corresponding to their spectral types, indicating that these lines are being filled in to some degree. Other emission lines present in the spectra are forbidden lines of [O i] and [S ii]. These lines are thought to be formed in extended, low-density, collisionally excited gas (Finkenzeller 1985). We find that 31 stars out of a total of 63 (late F stars are not included) exhibit Wk of [O i] k6300 in emission larger than the typical lowest emission we can detect in our the sample, 0.1 Å. The bottom panel of Figure 6 shows the distribution of Wk of [O i] k6300. For a subset of 17 stars, the [O i] k6363 line could be measured. The mean ratio of [O i] k6300/k6363 for HAeBe and continuum stars is 2.7, close to the optically thin ratio (Osterbrock 1989). Emission in multiplets of Fe ii, most conspicuously in multiplets 42 and 49, is found in 25 objects, 33% of the sample. When the emission is present, it is related to the Balmer emission. As shown in Figure 7, a correlation exists between the Wk of Fe ii k5169 and the Wk of H. The correlation coefficient is 0.74. In general, Fe ii (42) is observed in emission only if [O i] k6300 is also present, except for BD +46 3471, MacC H1, and LkH 218. However, Böhm & Catala (1994) detected emission at [O i] k6300 in two of these stars: in BD +46 3471 with a Wk ¼ 0:1 Å, and in LkH 218 with a Wk ¼ 0:2 Å. Still, no clear correlation between the strengths of the lines is found; the correlation coefficient between the equivalent widths is 0.43. A fraction of stars show emission in He i lines. Comparison of Tables 6 and 8 shows that this emission is more frequent among the continuum stars than in normal HAeBe stars (57% for the continuum stars, 18% for HAeBe stars). TABLE 8 Nonphotospheric Spectral Features 1694 HBC Name 1........ 3........ 7........ 78...... 94...... 154.... 160.... 169.... 170.... 192.... 193.... 196.... 197.... 201.... 202.... 217.... 219.... 222.... 231.... 273.... 281.... 282.... 284.... 293.... 297.... 305.... 308.... 309.... 310.... 313.... 314.... 321.... 324.... 325.... 329.... 334.... 348.... 350.... 373.... 430.... 432.... 436.... 442.... 451.... 460.... 464.... MacC H12 V633 Cas LkH 201 AB Aur HK Ori T Ori PQ Ori BF Ori RR Tau HD 250550 LkH 208 LkH 338 LkH 339 LkH 341 VY Mon W84 V590 Mon W108 V360 Mon KK Oph LkH 118 VV Ser AS 310 NW PX Vul V751 Cyg LkH 324 LkH 349 LkH 234 BD +46 3471 LkH 233 LkH 350 MacC H4 MC1 V376 Cas VX Cas RNO 6 IP Per XY Per EW V892 Tau UX Ori Par 102 RY Ori P1394 HD 245185 MV Ori CQ Tau H H [O i] [O i] [S ii] [S ii] [N ii] Ca ii Fe ii Fe ii Fe ii Fe ii Fe ii Fe ii Fe ii He i He i m=1 m=1 m=1 m=1 m =2 m =2 m=1 m=1 m = 42 m = 42 m = 42 m = 49 m = 49 m = 49 m = 48, 49 m = 11 m = 46 Typesa k0 = 6563 k0 = 4861 k0 = 6300 k0 = 6363 k0 = 6717 k0 = 6731 k0 = 6583 k0 = 3934 k0 = 4924 k0 = 5018 k0 = 5169 k0 = 5198 k0 = 5235 k0 = 5276 k0 = 5317 k0 = 5876 k0 = 6678 n h u h n h u n h h h h h u n f h f f n u h h h u h f h h h u n h n h h h h h h f f f h f h 31.5 56.2 41.6 28.2 49.0 21.0 ... 6.7 25.7 24.8b 4.9 51.0 19.4 23.5 28.0 10.3 47.3 1.2 20.7 59.4 19.1 61.1 7.7 6.4 7.9 15.3 0.3b 68.9 18.6 20.5 29.0 24.0 16.9 22.0 19.2 0.2 21.4 <4.7c 17.8 2.3 6.6 7.0 1.3b 21.2 ... 6.2 2.3 4.0 3.5 1.6 4.1 ... ... ... ... 3.3b ... 3.4 0.80 3.6 0.6b 0.3 0.5 ... 1.5 1.1 2.4 1.3 ... ... 1.8 0.5 ... 5.2 0.6 1.1 3.2 4.3 ... ... ... ... 0.9 ... ... ... ... ... ... ... ... ... 3.9 1.9 ... 0.1 1.2 ... ... ... 0.3 0.1 ... 0.3 ... ... 1.4 ... 1.4 ... ... 2.2 ... 0.6 ... ... ... ... ... 0.9 ... 0.7 ... 0.7 0.6 1.7 0.2 ... ... ... 0.3 ... ... ... ... ... ... ... 1.4 0.6 ... ... 0.5 ... ... ... ... ... ... ... ... ... 0.4 ... 0.5 ... ... 0.5 ... ... ... ... ... ... ... 0.3 ... 0.3 ... ... 0.2 0.8 ... ... ... ... ... ... ... ... ... ... ... ... 1.6 0.3 ... ... ... ... ... ... ... ... ... ... ... ... 0.2 ... ... ... ... 0.2 ... ... 0.6 ... ... ... ... ... ... 0.4 ... ... 0.2 0.9 ... 0.2 ... ... ... ... ... ... ... ... ... ... 2.5 0.4 ... ... ... ... ... ... ... ... ... ... ... ... 0.4 ... ... ... ... 0.3 ... ... 0.9 ... ... ... ... ... ... 0.5 ... ... 0.3 1.1 ... 0.2 ... ... ... ... ... ... ... ... ... ... 2.2 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 3.7 ... ... ... ... ... ... 0.5 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.6 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 1.3 ... 0.3 ... ... ... +1.4 +0.6 0.2 +0.5 0.30 ... ... 0.7 ... ... ... ... ... ... +0.2 ... ... ... ... ... 1.2 0.1 ... ... ... +0.6 ... ... ... ... +1.0 ... +0.8 ... ... ... ... ... ... ... 1.7 0.4 0.4 0.6 ... ... +1.7 +0.7 0.4 +0.5 0.75 ... ... ... ... ... ... ... 0.4 ... +0.2 ... ... ... ... ... 1.8 0.2 ... ... ... +0.6 ... ... ... ... +1.2 ... +0.8 ... ... ... ... ... ... ... 1.8 0.5 0.4 1.7 ... ... +2.4 +1.0 0.5 +0.9 0.53 ... ... 1.2 ... 0.4 ... ... 1.4 ... +0.3 ... ... ... ... ... 2.1 0.2 ... ... ... +1.2 ... ... ... ... +1.8 ... +1.1 ... ... ... ... ... ... ... 0.5 0 ... ... ... ... +0.3 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.7 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.4 0.3 ... ... ... ... +0.7 ... ... ... ... ... ... 0.4 ... ... ... ... ... ... ... ... ... ... ... ... 0.7 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.8 0.2 ... 0.4 ... ... +1.0 ... ... ... ... ... ... 0.6 ... ... ... ... ... ... ... ... ... ... ... ... 1.0 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.8 0.4 ... 0.6 ... ... +0.7 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 1.3 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.9 ... ... 0.6 ... ... ... +0.9 ... ... ... ... 0.3 ... ... 0.4 ... ... ... ... 1.3 ... ... ... 0.2 ... ... ... ... 0.3 ... 0.6 ... ... ... ... 0.5 ... ... ... ... ... ... 0.2 ... ... ... ... ... 0.30 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.6 ... ... ... 0.70 ... ... ... ... ... ... 0.7 ... ... ... ... ... ... ... ... ... ... ... ... ... ... TABLE 8—Continued 1695 HBC Name 482.... 492.... 493.... 518.... 528.... 529.... 531.... 535.... 548.... 551.... 686.... 689.... 690.... 705.... 716.... 717.... 726.... 730.... 734.... 735.... 736.... 742.... BN Ori BD +26 887 V350 Ori RNO 63 LkH 215 HD 259431 VSB 2 W121 LkH 218 LkH 222 WW Vul V1685 Cyg V1686 Cyg LkH 147 V1493 Cyg LkH 168 HD 200775 BD +65 1637 BH Cep BO Cep SV Cep MacC H1 H H [O i] [O i] [S ii] [S ii] [N ii] Ca ii Fe ii Fe ii Fe ii Fe ii Fe ii Fe ii Fe ii He i He i m=1 m=1 m=1 m=1 m =2 m =2 m=1 m=1 m = 42 m = 42 m = 42 m = 49 m = 49 m = 49 m = 48, 49 m = 11 m = 46 Typesa k0 = 6563 k0 = 4861 k0 = 6300 k0 = 6363 k0 = 6717 k0 = 6731 k0 = 6583 k0 = 3934 k0 = 4924 k0 = 5018 k0 = 5169 k0 = 5198 k0 = 5235 k0 = 5276 k0 = 5317 k0 = 5876 k0 = 6678 u h h f h h f f h h h h f h n n h h h h h n 1.3b 3.6 29.9 2.3b 25.7 57.5 <0 <0c 32.3 54.5 14.4 108 3.6 26.9 9.5 19.0 59.3 28.0 <6.2c <2.3c 12.1 34.2b ... ... ... ... 0.3 2.9 ... ... 1.4 2.9 ... 8.9 ... 1.1 ... ... 2.7 2.2 ... ... ... 4.1b ... ... 0.4 ... ... 0.5 ... ... ... 0.3 ... 0.9 ... ... 0.2 ... 0.1 ... ... ... ... ... ... ... 0.2 ... ... 0.2 ... ... ... 0.2 ... 0.3 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 0.8 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... +0.3 ... +0.7 0.3 ... ... 0.5 0.5 +0.6 0.3 ... ... +0.7 ... ... ... ... ... +0.4 0.6 ... ... +0.5 +0.3 0.5 ... ... 0.6 0.6 +0.6 0.8 ... ... +0.7 ... 0.1 0.5 ... ... +0.5 1.0 ... ... +0.9 ... +0.3 0.9 ... ... 0.4 0.5 +0.7 1.22 ... ... +1.0 ... 0.7 0.7 ... ... +0.6 1.0 Notes.—Here ‘‘m’’ indicates the multiplet of the element and k0 is the wavelength of the feature in angstroms. a Types are ‘‘h’’ HAeBe star; ‘‘f’’ late F star; ‘‘n’’ no spectral type could be assigned to this star; and ‘‘u’’ stars with uncertain evolutionary status. b P Cygni profile. c Double-peaked profile. ... ... ... ... ... 0.2 ... ... ... ... ... 0.4 ... ... ... ... ... 0.3 ... ... ... 0.5 ... ... ... ... ... 0.4 ... ... ... ... ... 0.5 ... ... ... ... ... 0.4 ... ... ... 0.6 ... ... ... ... ... 0.5 ... ... ... ... ... 0.7 ... ... ... ... ... 0.4 ... ... ... 0.4 ... ... ... ... ... 0.6 ... ... ... ... ... 0.7 ... ... ... ... ... 0.5 ... ... ... 0.6 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... 1696 HERNÁNDEZ ET AL. Vol. 127 phenomena are related. Models of UX Ori objects invoke obscuration of the star by circumstellar material crossing the line of sight; this could be produced by large orbiting circumstellar clouds, infalling cometary bodies, or instabilities in a flared disk (Natta & Whitney 2000; Bertout 2000; Graham 1992; Grinin 1988). Recently Dullemond et al. (2003) proposed an alternative explanation for this phenomenon, in which the obscuring region is the inner rim of a truncated disk. In this model, the inner disk is puffed up as a result of the rim being much hotter than the rest of the disk, such that this inner rim shadows the outer vicinity of the disk (Dullemond et al. 2001). The same material occulting the star could be responsible for the anomalous absorption features. High-resolution studies indicate variations of Fe ii (42), going from a P Cygni profile to an absorption profile (Rodgers et al. 2002; Catala et al. 1993). 6. REDDENING TOWARD HAeBe STARS Fig. 6.—Distribution of Wk for H (top) and [O i] k6300 (bottom): HAeBe stars (solid histogram), late F stars (shaded histogram), and continuum stars (open histogram). Continuum stars exhibit the strongest emission at H . The late F stars do not show [O i] k6300 in emission. While emission features are more easily distinguished, the nonphotospheric absorption features, sometimes referred to as ‘‘shell’’ features, are more difficult to single out. However, in our classification scheme based on multiple indices, these shell features tend to stand out as yielding spectral types incongruent with those obtained from the rest of the indices. The more readily identified shell features at our spectral resolution correspond to lines of multiplet 42 of Fe ii. Other lines that may be affected by absorption external to the photosphere are the Na i D line and He i k5876; the last when observed at high dispersion often has blueshifted components (e.g., MWC 1080 He i profile in Hartmann, Kenyon, & Calvet 1993), arising in expanding material around the star. However, at our resolution individual shell components cannot be picked out. We have identified nonphotospheric absorption in the 42 multiplet of Fe ii in 11 HAeBe stars and the star BF Ori (without spectral type assigned in this work), with the Wk shown in Table 8. Among the objects with Fe ii nonphotospheric absorption, the stars BF Ori (HBC 169), RR Tau (HBC 170), LkH 208 (HBC 193), VV Ser (HBC 282), UX Ori (HBC 430), WW Vul (HBC 686), and SV Cep (HBC 736) are reported as belonging to the UX Ori class by Natta et al. (1997) and Grinin (1994). Moreover, the star XY Per (HBC 350) also shows photometric properties characteristic of UX Ori objects (Shevchenko et al. 1993; Chkhikvadze 2002), and V350 Ori (HBC 493) shows UX Ori type photometric and polarimetric properties (Yudin & Evans 1998). The stars V1493 Cyg (HBC 716) and MC 1 (HBC 324) do not have enough data to decide whether their photometric properties fall in the UX Ori group. On the other hand, the quasi-periodic light curve reported for LkH 215 (HBC 528; Shevchenko et al. 1993) and its small photometric variability range (Herbst & Shevchenko 1999) tend to argue against a UX Ori nature. Overall, 75% of the stars with anomalous Fe ii absorption have been reported as UX Ori, which strongly suggests that the two Pre–main-sequence stars have long been known to exhibit significant photometric variability across a wide wavelength range. In particular, 25% of all known HAeBe stars are reported to show strong variations in brightness and color, which can exceed 4 mag in some wavelength ranges like the Johnson V band (Finkenzeller & Mundt 1984; Herbst & Shevchenko 1999). The UX Ori objects fall within this group. (Rodgers et al. 2002; Natta et al. 2000; Natta & Whitney 2000). This behavior means that great care must be exercised when measuring a representative extinction (AV) toward these stars. Variability of the spectral type could further complicate reddening determinations. Ideally one would want to have multiepoch, simultaneous spectra and photometry for each star to estimate mean magnitudes and spectra, but this requires an observational effort that is seldom feasible (though some campaigns like EXPORT are aimed toward this multiepoch, multiwavelength approach). We adopt the spectral type found in this study as representative of the photosphere of the star and use an extensive Fig. 7.—Comparison between the Wk of H and the Fe ii k5169 line. Although there is significant scatter, a simple linear regression yields a correlation index of 0.74, suggesting a trend between these lines. No. 3, 2004 SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS photometric data set so the variability range can be readily assessed. The optical photometry we use consists of a large data set in the UBVR bands that Herbst & Shevchenko (1999) have amassed by monitoring a set of HAeBe stars since 1983; 69% of the stars that we have classified here have UBVR measurements in their work. We complemented the Herbst & Shevchenko (1999) data using measurements from de Winter et al. (2001), Miroshnichenko et al. (2001), Flaccomio et al. (1999), Fernandez (1995), Hillenbrand et al. (1992), Mendoza & Gomez (1980), Herbig & Bell (1988), and MacConnell (1968). We computed mean and median UBVR magnitudes for each star in our sample using these databases; these agreed within 0.1 mag or better. The most variable of the stars in our sample are the UX Ori type objects; for these, we find that the mean magnitudes correspond to the bright state when the star is probably seen without the occulting screen. Thus, in our analysis we adopt mean magnitudes and colors as representative of the brightness of the stellar photosphere. We calculated the color excesses EVB and EVR using intrinsic colors given for each spectral type by Kenyon & Hartmann (1995). With this information we obtained the values of the visual extinctions AV1 from EVB and AV 2 from EVR, with different values of the total-to-selective extinction RV (AV = RVEBV). We used the relations given in Cardelli et al. (1989) to calculate AðkÞ=AðV Þ for a specific RV . In Figure 8 we plot AV1 versus AV 2 for RV ¼ 3:1, consistent with the mean interstellar medium (left) and for RV ¼ 5 (right). The line with slope 1 is indicated in both cases. The values of the extinction derived from different colors agree if RV is significantly higher than 3.1; for RV ¼ 5 the correlation is AV 1 ¼ 1:002AV 2 0:136, while for RV ¼ 3:1 it is AV 1 ¼ 0:779AV 2 0:08. The significant points for determining these correlations are stars with high reddening, AV > 1:5, which constitute 73% of the sample. In order to determine if the reddening toward these stars is interstellar or circumstellar, we examine their distances, given in Table 9. We find that 85% of the highly reddened stars are located within 1 kpc from the Sun. Since for most lines of sight, the expected interstellar reddening is less than 1 mag for this range of distances (Fitzgerald 1968), we conclude that the high reddening is not interstellar in nature; in fact, this reddening is most likely produced by a combination of circumstellar material close to the star and material from the molecular clouds with which many of these objects are still associated. The high value 1697 of RV then strongly suggests that the circumstellar medium around HAeBe stars is dominated by grains larger in size than the average dust grain in the diffuse interstellar medium. Other authors have also arrived at a similar conclusion using smaller samples of these stars (Strom et al. 1972; Thé et al. 1981; Herbst et al. 1982; Sorrell 1990; Bibo et al. 1992; Gorti & Bhatt 1993; Waters & Waelkens 1998; Whittet et al. 2001). Studies of silicate features also show strong evidence of coagulation and an increase in average grain size (Meeus et al. 2002; Bouwman et al. 2001; Meeus et al. 2001) In Figure 9 we plot the ratio EVRC =EBV versus EVIC =EBV for a subsample of HAeBe stars with measured IC from the Van Vleck Observatory public ftp server,8 Fernandez (1995), Oudmaijer et al. (2001), and de Winter et al. (2001). We also plot the expected color ratios for different values of RV using the relations of Cardelli et al. (1989). For this, extinctions AðkÞ=AðV Þ were calculated at the effective wavelength of the filters, using tables of effective wavelength versus VIC kindly provided by M. Bessell. The two lines correspond to the minimum and maximum VIC of the sample. Again, the data suggest an extinction law with RV > 3:1. 7. LOCATION OF THE STARS IN THE H-R DIAGRAM Knowledge of the appropriate extinction law for HAeBe stars is important in order to derive luminosities for these stars, which in turn allow us to estimate their masses and evolutionary status. We calculated the stellar luminosity for 55 out of the 58 stars shown in Tables 2, 3, and 4. Three stars, P102, MV Ori, and RNO 63, did not have enough published photometric data to enable a reliable estimate of luminosity. We used the mean V magnitude given in Table 9, corrected for reddening with an AV obtained from mean colors (x 6), bolometric corrections from Kenyon & Hartmann (1995), and distances from the references cited in Table 9. The extinction correction was calculated by comparing the BV colors with intrinsic colors for the spectral type from Kenyon & Hartmann (1995), using both the standard interstellar extinction law RV ¼ 3:1 and RV ¼ 5:0. The effective temperature was determined using our spectral types and the calibration of Kenyon & Hartmann (1995). 8 See ftp://ftp.astro.wesleyan.edu/pub/ttauri. Fig. 8.—Comparison of reddening values AV determined from E(BV ) and E(VR), for RV ¼ 3:1 (left) and RV = 5.0 (right). The solid line represents the fit to the data, while the dashed line has slope unity. The best correlation is observed for RV ¼ 5:0. Error bars represent the propagated error from the spectral types. TABLE 9 Ages and Masses RV = 3.1 RV = 5.0 HBC Name V (Mag) Reference Distance (pc) Reference Teff (K) AV (mag) logL (L) Mass (M) Age (Myr) AV (mag) logL (L) Mass (M) Age (Myr) 3.......... 7.......... 78........ 154...... 160...... 170...... 192...... 193...... 196...... 197...... 201...... 217...... 219...... 222...... 231...... 281...... 282...... 284...... 293...... 297...... 305...... 308...... 309...... 310...... 313...... 314...... 324...... 329...... 334...... 348...... 350...... 373...... 430...... 436...... 442...... 451...... 464...... 482...... 492...... 493...... 528...... 529...... 531...... 535...... 548...... 551...... 686...... 689...... 690...... 705...... 726...... 730...... 734...... 735...... 736...... V633 Cas LkH 201 AB Aur T Ori PQ Tau RR Tau HD 250550 LkH 208 LkH 338 LkH 339 LkH 341 W84 V590 Mon W108 V360 Mon LkH 118 VV Ser AS 310 NW PX Vul V751 Cyg LkH 324 LkH 349 LkH 234 BD +46 3471 LkH 233 LkH 314 MC1 VX Cas RNO 6 IP Per XY Per EW V892 Tau UX Ori RY Tau P1394 HD 245185 CQ Tau BN Ori BD +26 887 V350 Ori LkH 215 HD 259431 VSB2 W121 LkH 218 LkH 222 WW Vul V1685 Cyg V1686 Cyg LkH 147 HD 200775 BD +65 1637 BH Cep BO Cep SV Cep 14.18 13.64 7.05 10.63 12.63 12.08 9.54 11.65 15.12 13.66 13.39 12.02 12.77 11.97 13.39 11.20 11.92 12.45 11.49 14.18 12.61 13.37 12.21 9.89 13.56 14.04 10.77 11.28 14.52 10.47 9.21 15.25 10.40 11.80 10.13 9.89 10.27 9.67 10.47 11.47 10.54 8.73 13.33 10.80 11.87 11.81 10.74 10.69 14.06 14.46 7.37 10.18 11.16 11.60 10.98 1 1 1 1 2 1 1 1 2 1 1 3 1 2 3 1 1 1 4 1 4 4 1 1 1 1 5 1 2 6 1 1 1 7 2 1 1 1 2 7 1 1 2 2 1 1 1 1 1 1 1 1 1 1 1 600 850 144 460 460 800 700 1000 830 830 800 910 800 910 758 1950 440 2500 420 700 780 750 1000 900 880 400 850 760 1600 350 120 160 460 460 460 400 130 460 2000 460 800 800 910 910 1150 1150 550 980 980 800 429 1250 450 400 400 8 9 10 8 8 8 8 8 11 11 12 13 8 13 14 15 16 8 12 17 18 19 16 16 8 20 9 8 8 21 10 8 8 8 8 8 22 8 23 8 21 8 13 13 8 12 16 22 22 24 10 22 20 20 20 4.03 4.32 3.97 3.98 3.83 3.99 4.04 3.91 4.05 3.97 3.83 3.80 4.11 3.80 3.76 4.43 4.14 4.40 3.83 3.99 4.09 3.79 4.12 3.99 3.93 4.08 3.90 3.99 4.26 3.92 3.92 4.05 3.94 3.80 3.79 3.97 3.83 3.82 3.92 3.96 4.14 4.15 3.77 3.77 3.98 4.09 3.94 4.27 3.79 4.32 4.27 4.22 3.81 3.82 4.00 3.2 4.4 0.3 1.6 0.6 2.0 0.4 0.8 3.3 2.6 1.8 0.2 0.8 0.3 0.6 3.5 3.4 4.1 1.4 0.9 3.7 3.2 3.1 0.3 2.3 6.3 0.4 1.0 2.3 0.6 1.1 4.8 0.9 1.2 0.3 0.2 1.2 0.3 0.8 1.3 2.0 1.2 0.4 0.0 1.3 1.2 1.0 3.0 2.5 5.3 1.8 1.8 0.7 0.5 1.3 1.28 2.96 1.63 1.73 0.48 1.81 2.20 1.59 1.26 1.43 1.13 1.15 1.32 1.21 0.64 4.55 2.23 4.43 1.15 0.45 2.48 1.65 2.67 2.12 1.34 2.31 1.67 1.70 2.18 1.07 0.86 0.41 1.49 1.05 1.35 1.31 0.55 1.51 2.69 1.24 2.75 3.19 0.74 1.58 1.91 2.12 1.56 3.58 1.32 2.94 3.73 3.40 1.06 0.72 1.38 a a a a 5.1 7.1 0.5 2.6 1.0 3.2 0.7 1.2 5.3 4.3 3.0 0.3 1.3 0.5 1.0 5.7 5.4 6.6 2.2 1.5 6.0 5.2 5.0 0.5 3.7 10.1 0.7 1.7 3.8 0.9 1.7 7.8 1.4 1.9 0.4 0.3 2.0 0.4 1.3 2.1 3.2 2.0 0.7 0.0 2.1 1.9 1.6 4.9 4.0 8.6 3.0 2.9 1.1 0.8 2.1 2.06 4.04 1.71 2.13 0.64 2.29 2.31 1.78 2.06 2.08 1.58 1.19 1.52 1.28 0.80 5.42 3.06 5.43 1.48 0.68 3.39 2.43 3.43 2.19 1.90 3.85 1.78 1.95 2.75 1.21 1.12 1.60 1.71 1.34 1.41 1.36 0.85 1.58 2.89 1.55 3.24 3.50 0.85 1.58 2.23 2.41 1.81 4.33 1.93 4.24 4.17 3.83 1.23 0.85 1.69 3.2 11.0 2.6 3.5 1.5 3.9 3.9 2.8 3.2 3.4 2.6 2.1 2.59 0.09 3.94 1.68 25.06 1.30 1.44 2.87 2.68 1.83 3.13 5.38 a 2.5 2.6 1.5 2.7 3.6 2.4 a 2.3 1.8 2.0 4.33 3.89 b 3.69 1.80 4.21 a 7.00 7.90 6.27 a a 2.1 1.5 16.2 3.8 14.5 1.9 5.18 12.49 0.04 a 4.3 3.0 4.8 3.5 2.1 3.9 2.5 2.7 a 1.9 a a 2.2 1.8 2.3 2.1 1.5 2.5 5.2 2.0 4.8 6.6 1.5 3.0 3.0 3.4 2.4 8.8 2.3 a 8.9 7.0 1.8 1.5 2.3 b 0.02 7.53 a 1.25 1.89 0.83 1.81 7.16 1.51 3.75 3.96 a 14.36 a a 5.74 7.62 3.80 9.95 b 3.51 0.45 b 0.80 0.32 11.74 1.88 2.80 2.78 4.48 0.13 3.69 a 0.13 0.29 8.22 14.41 b a 2.2 1.7 42.6 5.8 43.5 2.4 a 7.9 5.4 7.9 3.7 2.9 11.6 2.8 3.0 5.0 2.0 1.9 2.8 2.5 2.3 2.4 2.2 1.6 2.7 6.2 2.3 6.6 9.7 1.7 3.0 3.8 4.1 2.9 64.1 3.8 12.9 12.5 25.0 2.1 1.6 2.7 a 4.27 8.81 0.01 0.46 0.01 3.70 a 0.21 0.30 0.22 1.52 2.61 0.05 2.83 2.73 b 7.98 10.92 b 4.02 3.86 3.46 7.97 13.27 2.93 0.28 5.64 0.30 0.09 9.25 1.88 1.40 1.38 2.94 0.11 0.95 0.06 0.06 0.14 5.47 12.48 3.96 The star falls below the ZAMS. The star falls on the ZAMS. References.—(1) Herbst & Shevchenko 1999; (2) Herbig & Bell 1988; (3) Flaccomio et al. 1999; (4) Fernandez 1995; (5) MacConnell 1968; (6) Miroshnichenko et al. 2001; (7) de Winter et al. 2001; (8) Testi et al. 1998; (9) Yonekura et al. 1997; (10) Bertout, Robichon, & Arenou 1999; (11) Herbst & Racine 1976; (12) Herbst et al. 1982; (13) Neri, Chavarria-K., & de Lara 1993; (14) Park et al. 2000; (15) Kozok 1985; (16) Pirzkal et al. 1997; (17) Chavarria et al. 1989; (18) Chavarria et al. 1983; (19) Hessman et al. 1995; (20) Kun 1998; (21) Hillenbrand et al. 1992; (22) van den Ancker et al. 1998; (23) Kawamura et al. 1998; (24) Natta et al. 2001. b SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS 1699 Figure 10 shows the location in the H-R diagram of these stars, for the two values of RV used. We also show the evolutionary tracks and isochrones of Siess et al. (2000) and Bernasconi (1996). We derive masses and ages for the sample by linear interpolation in these tracks. The derived values are given in Table 9 for each value of RV. Clearly the value of RV makes a significant difference in the H-R diagram positions of the sample. When using RV ¼ 3:1, 12 objects fall below the zero main-sequence age (ZAMS) and many close to it. In contrast, for RV ¼ 5:0 only two objects appear below the ZAMS, V751 Cyg (HBC 297) and V590 Mon (HBC 219). As discussed in see x 4.1, V751 Cyg probably is a cataclysmic variable. Testi et al. (1998) published for V590 Mon a BV color that is 1 mag larger than any value published before. With their BV measurement, V590 Mon falls well above the zero mainsequence age (ZAMS). However, a large variation in BV is difficult to reconcile with the fact that this star shows little photometric variability with a quasi-periodic light curve (Herbst & Shevchenko 1999). 8. SUMMARY AND CONCLUSIONS Fig. 9.—Ratio of color excesses EVRc /EBV versus ratio of color excesses EVIC =EBV . The dotted lines indicate the locus of predicted color excess ratios for values of RV from 2 to 8, calculated using the Cardelli et al. (1989) reddening law. The extinctions were calculated at the effective wavelength of the filters, for the minimum (top dotted line) and maximum (bottom dotted line) VIC of our sample (see text). Filled circles are objects with more than five photometric measurements. The asterisks correspond to stars with only one or two measurements in each filter; for these objects variability could be affecting our estimate of color excesses. We have applied a consistent spectral classification scheme aimed at early-type PMS stars. Our method relies on a large number of reddening-independent indices covering a wide wavelength range, from 3900 to 7000 Å. Our scheme is designed to avoid contamination by nonphotospheric contributions to absorption features normally used for spectral typing. We were able to determine spectral types for 58 objects out of a total sample of 75 objects from the (Herbig & Bell 1988) catalog, with an average uncertainty of 2.5 subclasses. Seven stars with spectral types have an uncertain evolutionary status. We could not derive spectral types for a Fig. 10.—Location of the stars in the H-R diagram. Left: Luminosities obtained using the standard value of total-to-selective extinction for the interstellar medium, RV ¼ 3:1. Right: Results using the larger value RV ¼ 5:0. When using RV ¼ 3:1, many stars fall on or below the ZAMS. A value of RV ¼ 5 tends to yield higher luminosities, moving the stars upward in the H-R diagram, hence making them younger. Then almost all stars fall above the ZAMS, which is more consistent with their pre–main-sequence nature. We show the evolutionary tracks (solid lines) and isochrones (dashed lines). Tracks represent, from top to bottom, 25, 15, and 9 M (Bernasconi 1996) and 6, 3, and 1 M (Siess et al. 2000). The isochrones from Siess et al. (2000) are, from top to bottom, 0.1, 1, 10, and 100 Myr (which we take as the ZAMS). Luminosity errors represent the propagated error from the spectral type. 1700 HERNÁNDEZ ET AL. subsample of 11 stars. In seven of these stars, no absorption features where apparent in the spectrum, which is dominated by emission lines; these are continuum stars. The contamination by nonstellar spectral features was too strong for the rest of this subsample, precluding spectral typing. Finally, only approximate spectral types are given for six highly reddened stars, for which the blue end of the spectrum was too faint to properly apply our classification scheme. By definition, all the stars of the sample show H in emission. However, at our resolution only 53% show H in emission as well, and only 15% show additional emission in H. Nonetheless, 95% and 56% of the HAeBe stars show filling-in of H and H, respectively. Almost half of the HAeBe stars classified in this work (excluding stars F7 and later) exhibit the forbidden line [O i] k6300 in emission (similar to reports by Corcoran & Ray 1997 and Böhm & Catala 1994). This feature is indicative of the presence of winds, outflows, or jets. A third of the sample exhibits emission in multiplets of Fe ii, particularly multiplet 42, and the strength of this emission is correlated with that of H . Emission in multiplets of Fe ii only appears if [O i] k6300 is present, although their strengths do not seem to correlate. A subset of 11 HAeBe stars, 28% of the HAeBe sample, shows lines of multiplet 42 of Fe ii with abnormally strong absorption. Of these stars, 75% have been confirmed as UX Ori objects, strongly suggesting that the anomalous Fe ii absorption is produced by the same mechanism that results in the UX Ori phenomenon. In fact, if we assume that the sample of HAeBe stars was complete, the number of objects with anomalous Fe ii absorption would be consistent with the expected number of high-inclination (>75 ) systems, which is a condition for the UX Ori phenomenon to occur (Natta & Whitney 2000). We have used published photometric data together with our derived spectral types to estimate the reddening law that best characterizes the class of HAeBe. We find that a reddening law with a high value of RV , 5.0, yields a much better agreement between values of the extinction AV obtained from different Vol. 127 colors than the standard reddening law. Since 85% of the stars with large values of AV are located within 1 kpc from the Sun, the high extinction values are probably not due to interstellar reddening. Rather, the stars must be mostly extincted by their circumstellar environments. Thus, the high value of RV that characterizes the reddening law toward the intermediate mass PMS stars indicates that dust has grown with respect to the typical grain size in the interstellar medium. Using reddening values determined for different extinction laws, we locate the stars in the H-R diagram. The position of the stars depends critically on the value chosen for RV, hence affecting estimates of masses and most particularly of ages. With the most appropriate value of RV ¼ 5:0, objects appear systematically younger and brighter relative to their positions calculated with the standard law. In particular, the majority of the spectroscopically selected, young, bona fide HAeBe stars consistently falls above the ZAMS. Additional information about the spectra, H emission, UBVRIJHK magnitudes, observed emission lines, finding charts, and optical and near infrared spectral energy distribution for each object analyzed in this work are reported on the World Wide Web.9 We thank Michael Bessell for sending us the color-dependence of the filter effective wavelengths, Bruno Merı́n for sending us the multiepoch spectra of the star V1686 Cyg, George Herbig for insightful conversations, and G. Meeus, the referee, for his careful reading of the manuscript and his detailed and useful comments and suggestions. We also thank Susan Tokarz of the SAO Telescope Data Center for carrying out the data reduction and Michael Calkins for obtaining some of the spectra. This work was supported in part by NASA grants NAG5-9670 and NAG10545, NSF grant AST 9987367 and grant S1-2001001144 of FONACIT, Venezuela. 9 See http://cfa-www.harvard.edu/youngstars/jhernand/haebe/principal.html. REFERENCES Arce, H. G., & Goodman, A. A. 2002, ApJ, 575, 911 de Winter, D., van den Ancker, M. E., Maira, A., Thé, P. S., Djie, H. R. E. T. A., Asselin, L., Menard, F., Bastien, P., Monin, J., & Rouan, D. 1996, ApJ, Redondo, I., Eiroa, C., & Molster, F. J. 2001, A&A, 380, 609 472, 349 Downes, R., Hoard, D. W., Szkody, P., & Wachter, S. 1995, AJ, 110, 1824 Bailey, J. 1998, MNRAS, 301, 161 Dullemond, C. P., Dominik, C., & Natta, A. 2001, ApJ, 560, 957 Bernasconi, P. A. 1996, A&AS, 120, 57 Dullemond, C. P., van den Ancker, M. E., Acke, B., & van Boekel, R. 2003, Bertout, C. 2000, A&A, 363, 984 ApJ, 594, L47 Bertout, C., Robichon, N., & Arenou, F. 1999, A&A, 352, 574 Echevarrı́a, J., Costero, R., Tovmassian, G., Zharikov, S., Pineda, L., & Michel, R. Bibo, E. A., Thé, P. S., & Dawanas, D. N. 1992, A&A, 260, 293 2002, Rev. Mexicana Astron. Astrofis. Ser. Conf., 12, 86 Böhm, T., & Catala, C. 1994, A&A, 290, 167 Fabricant, D., Cheimets, P., Caldwell, N., & Geary, J. 1998, PASP, 110, 79 ———. 1995, A&A, 301, 155 Fernandez, M. 1995, A&AS, 113, 473 Bouwman, J., Meeus, G., de Koter, A., Hony, S., Dominik, C., & Waters, Fernandez, M., Ortiz, E., Eiroa, C., & Miranda, L. F. 1995, A&AS, 114, 439 L. B. F. M. 2001, A&A, 375, 950 Finkenzeller, U. 1985, A&A, 151, 340 Briceno, C., Hartmann, L. W., Stauffer, J. R., Gagne, M., Stern, R. A., & Finkenzeller, U., & Jankovics, I. 1984, A&AS, 57, 285 Caillault, J. 1997, AJ, 113, 740 Finkenzeller, U., & Mundt, R. 1984, A&AS, 55, 109 Buscombe, W. 2001, VizieR On-line Data Catalog, III/222 Fitzgerald, M. P. 1968, AJ, 73, 983 Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245 Flaccomio, E., Micela, G., Sciortino, S., Favata, F., Corbally, C., & Tomaney, A. Casey, S. C., & Harper, D. A. 1990, ApJ, 362, 663 1999, A&A, 345, 521 Catala, C., Böhm, T., Donati, J.-F., & Semel, M. 1993, A&A, 278, 187 Garcia, B. 1989, Bull. Inf. Centre Donnees Stellaires, 36, 27 Chavarria, C., de Lara, E., Finkenzeller, U., Appenzeller, I., & Cardona, O. Gomez, M., Kenyon, S. J., & Whitney, B. A. 1997, AJ, 114, 265 1983, A&A, 118, 189 Gorti, U., & Bhatt, H. C. 1993, A&A, 270, 426 Chavarria-K., C., Terranegra, L., Alcala, J. M., & Neri, L. 1989, Rev. Mexicana Graham, J. A. 1992, PASP, 104, 479 Astron. Astrofis., 18, 178 Gray, D. F. 1992, The Observation and Analysis of Stellar Photospheres Chkhikvadze, I. N. 2002, Astrophysics, 45, 150 (2d ed.; Cambridge: Cambridge Univ. Press), 81 Cohen, M., & Kuhi, L. V. 1979, ApJS, 41, 743 Gray, R. O., Napier, M. G., & Winkler, L. I. 2001, AJ, 121, 2148 Coluzzi, R. 1999, VizieR On-line Data Catalog, VI/71A Grinin, V. P. 1988, Soviet Astron. Lett., 14, 27 Corcoran, M., & Ray, T. 1997, A&A, 321, 189 ———. 1994, ASP Conf. Ser. 62, The Nature and Evolutionary Status of Corporon, P., & Lagrange, A.-M. 1999, A&AS, 136, 429 Herbig Ae/Be Stars, ed. P. S. Thé, M. R. Pérez, & E. P. J. van den Heuvel Davies, J. K., Evans, A., Bode, M. F., & Whittet, D. C. B. 1990, MNRAS, (San Francisco: ASP), 63 247, 517 Hajjar, R., & Bastien, P. 2000, ApJ, 531, 494 No. 3, 2004 SPECTRAL ANALYSIS OF HERBIG Ae/Be STARS Hamann, F., & Persson, S. E. 1992, ApJS, 82, 285 Hartmann, L., Calvet, N., Gullbring, E., & D’Alessio, P. 1998, ApJ, 495, 385 Hartmann, L., Kenyon, S. J., & Calvet, N. 1993, ApJ, 407, 219 Herbig, G. H. 1960, ApJS, 4, 337 Herbig, G. H., & Bell, K. 1988, Lick Obs. Bull., 1111 Herbst, W., & Racine, R. 1976, AJ, 81, 840 Herbst, W., & Shevchenko, V. S. 1999, AJ, 118, 1043 Herbst, W., Warner, J. W., Miller, D. P., & Herzog, A. 1982, AJ, 87, 98 Hessman, F. V., Beckwith, S. V. W., Bender, R., Eisloeffel, J., Goetz, W., & Guenther, E. 1995, A&A, 299, 464 Hillenbrand, L. A. 1995, Ph.D. thesis, Univ. Massachusetts Hillenbrand, L. A., Strom, S. E., Vrba, F. J., & Keene, J. 1992, ApJ, 397, 613 Jaschek, M. 1978, Bull. Inf. Centre Donnees Stellaires, 15, 121 Kawamura, A., Onishi, T., Yonekura, Y., Dobashi, K., Mizuno, A., Ogawa, H., & Fukui, Y. 1998, ApJS, 117, 387 Keenan, P. C., & Barnbaum, C. 1999, ApJ, 518, 859 Kenyon, S. J., & Hartmann, L. 1995, ApJS, 101, 117 Kozok, J. R. 1985, A&AS, 62, 7 Kun, M. 1998, ApJS, 115, 59 Leinert, C., Richichi, A., & Haas, M. 1997, A&A, 318, 472 Lorenzetti, D., Saraceno, P., & Strafella, F. 1983, ApJ, 264, 554 MacConnell, D. J. 1968, ApJS, 16, 275 Magakian, T. Y., & Movsesian, T. A. 2001, Astrophysics, 44, 419 Malfait, K., Bogaert, E., & Walkens, C. 1998, A&A, 331, 211 Mannings, V., & Sargent, A. I. 1997, ApJ, 490, 792 ———. 2000, ApJ, 529, 391 Marconi, M., Ripepi, V., Bernabei, S., Palla, F., Alcalà, J. M., Covino, E., & Terranegra, L. 2001, A&A, 372, L21 Meeus, G., Bouwman, J., Dominik, C., Waters, L. B. F. M., & de Koter, A. 2002, A&A, 392, 1039 Meeus, G., Waters, L. B. F. M., Bouwman, J., van den Ancker, M. E., Waelkens, C., & Malfait, K. 2001, A&A, 365, 476 Mendoza, V. E. E., & Gomez, T. 1980, MNRAS, 190, 623 Miroshnichenko, A. S., Bjorkman, K. S., Chentsov, E. L., Klochkova, V. G., Gray, R. O., Garcı́a-Lario, P., & Perea Calderón, J. V. 2001, A&A, 377, 854 Miroshnichenko, A. S., et al. 2003, A&A, 408, 305 Mora, A., et al. 2001, A&A, 378, 116 Morgan, W. W., Keenan, P. C., & Kellman, E. 1943, An Atlas of Spectra Classification (Chicago: Univ. Chicago Press) Muzerolle, J., Calvet, N., & Hartmann, L. 2001, ApJ, 550, 944 Muzerolle, J., Calvet, N., Hartmann, L., Briceño, C., Hillenbrand, L. A., & Hernández J. 2004, in preparation Natta, A., Grinin, V., & Mannings, V. 2000, in Protostars and Planets IV, ed. V. Manning, A. P. Boss, & S. S. Russell (Tucson: Univ. Arizona Press), 559 Natta, A., Grinin, V. P., Mannings, V., & Ungerechts, H. 1997, ApJ, 491, 885 Natta, A., Prusti, T., Neri, R., Wooden, D., Grinin, V. P., & Mannings, V. 2001, A&A, 371, 186 1701 Natta, A., & Whitney, B. A. 2000, A&A, 364, 633 Neri, L. J., Chavarria-K., C., & de Lara, E. 1993, A&AS, 102, 201 Osterbrock, D. E. 1989, Astrophysics of Gaseus Nebulae and Active Galactic Nuclei (Mill Valley: Univ. Science Books) Oudmaijer, R. D., et al. 2001, A&A, 379, 564 Palla, F., & Stahler, S. W. 1991, ApJ, 375, 288 Park, B., Sung, H., Bessell, M. S., & Kang, Y. H. 2000, AJ, 120, 894 Parsamian, E. S., Gasparian, K. G., & Ohanian, G. B. 1996, Astrophysics, 39, 121 Pirzkal, N., Spillar, E. J., & Dyck, H. M. 1997, ApJ, 481, 392 Pritchet, C., & van den Bergh, S. 1977, ApJS, 34, 101 Reid, I. N., Hawley, S. L., & Gizis, J. E. 1995, AJ, 110, 1838 Robinson, E. L. 1973, ApJ, 180, 121 Rodgers, B., Wooden, D. H., Grinin, V., Shakhovsky, D., & Natta, A. 2002, ApJ, 564, 405 Shevchenko, V. S., Ezhkova, O., Tjin A Djie, H. R. E., van den Ancker, M. E., Blondel, P. F. C., & de Winter, D. 1997, A&AS, 124, 33 Shevchenko, V. S., Grankin, K. N., Ibragimov, M. A., Mel’Nikov, S. Y., & Yakubov, S. D. 1993, Ap&SS, 202, 121 Siess, L., Dufour, E., & Forestini, M. 2000, A&A, 358, 593 Skumanich, A. 1972, ApJ, 171, 565 Sorelli, C., Grinin, V. P., & Natta, A. 1996, A&A, 309, 155 Sorrell, W. H. 1990, ApJ, 361, 150 Stock, J., & Stock, J. M. 1999, Rev. Mexicana Astron. Astrofis., 35, 143 Strom, K. M., Wilkin, F. P., Strom, S. E., & Seaman, R. L. 1989, AJ, 98, 1444 Strom, S. E. 1983, Rev. Mexicana Astron. Astrofis., 7, 201 Strom, S. E., Strom, K. M., Yost, J., Carrasco, L., & Grasdalen, G. 1972, ApJ, 173, 353 Terranegra, L., Chavarria-K., C., Diaz, S., & Gonzalez-Patino, D. 1994, A&AS, 104, 557 Testi, L., Palla, F., & Natta, A. 1998, A&AS, 133, 81 Thé, P. S., de Winter, D., & Perez, M. R. 1994, A&AS, 104, 315 Thé, P. S., et al. 1981, A&AS, 44, 451 van den Ancker, M. E., de Winter, D., & Tjin A Djie, H. R. E. 1998, A&A, 330, 145 van den Ancker, M. E., Thé, P. S., Feinstein, A., Vazquez, R. A., de Winter, D., & Perez, M. R. 1997, A&AS, 123, 63 Waters, L. B. F. M., & Waelkens, C. 1998, ARA&A, 36, 233 Whitney, B. A., & Hartmann, L. 1993, ApJ, 402, 605 Whittet, D. C. B., Gerakines, P. A., Hough, J. H., & Shenoy, S. S. 2001, ApJ, 547, 872 Wu, Y., Huang, M., & He, J. 1996, A&AS, 115, 283 Yonekura, Y., Dobashi, K., Mizuno, A., Ogawa, H., & Fukui, Y. 1997, ApJS, 110, 21 Yudin, R. V., & Evans, A. 1998, A&AS, 131, 401