Download a0041_Spectral Analysis and Classification of Herbig A

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].
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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.
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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.
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