A&A 365, 90-109 (2001)
DOI: 10.1051/0004-6361:20000018
D. F. M. Folha1 - J. P. Emerson2
Send offprint request: D. F. M. Folha
1 - Centro de Astrofísica da Universidade do Porto, Rua
das Estrelas, 4150-762 Porto, Portugal
2 -
Department of Physics, Queen Mary, University of London,
Mile End Road, London E1 4NS, UK
Received 5 July 2000 / Accepted 6 October 2000
Abstract
From a sample of 50 T Tauri stars, mostly from the Taurus-Auriga
complex, Pa
line profiles were obtained for 49 of the
stars and Br
profiles for 37 of the stars. Emission
at Pa
was observed for 42 stars and emission at
Br
was found for 30 stars. The most conspicuous
features in the line profiles is the almost complete absence of
blueshifted absorption components and the relatively high frequency of
inverse P Cygni profiles (IPC). At Pa
,
34% of the
profiles are IPC while at Br
20% are IPC. The
redshifted absorption features indicate infall at velocities of about
,
compatible with free fall from a few radii
out. In general, line profiles are broad centrally peaked with
slightly blueshifted line peaks. Existing wind and accretion models
fail, in quantitative terms, to explain the shape of the observed
profiles. Magnetospheric accretion models, being the currently
preferred ones, produce lines too narrow (by
FWHM), wings extending to velocities too low (by at least
)
and with maximum normalized
intensities too high by factors of a few. A qualitative agreement
between some of the accretion model predicted profiles and some
observations hint that emission in these lines might, at least
partially, arise from infalling material. Definite claims regarding
the origin of the emission in these lines cannot be made until models
match observations much better than they currently do.
Key words: line: profiles - circumstellar matter - stars: formation - stars: pre-main sequence - infrared: general - infrared: stars
Author for correspondance: j.p.emerson@qmw.ac.uk
Hydrogen lines in emission are the hallmark of Classical T Tauri stars (CTTS). These lines tend to reflect the dynamics of the region where they are formed. Clear examples are normal or inverse P Cygni profiles found in the Balmer lines of many CTTS. Traditionally, the strong hydrogen emission was interpreted in terms of mass loss (Lago 1979; De Campli 1981; Hartmann et al. 1982; Natta et al. 1988; Hartmann & Kenyon 1990). The last decade saw a shift in interpretation towards mass accretion (Calvet & Hartmann 1992; Hartmann et al. 1994; Edwards et al. 1994; Muzerolle et al. 1998a; Muzerolle et al. 1998b), following the magnetospheric accretion model proposed by Camenzind (1990), Königl (1991) and Shu et al. (1994).
Considerable effort has been made in trying to model line profiles, especially
those of H
,
both in wind scenarios (Hartmann et al.
1990; Calvet et al. 1992; Grinin & Mitskevich 1991;
Mitskevich et al. 1993; Pedrosa 1996) and accretion
scenarios (Bertout 1977; Bertout 1979;
Bastian 1982; Calvet & Hartmann 1992;
Hartmann et al. 1994; Muzerolle et al. 1998a).
This has been accompanied by a similar effort on the observational
side, aiming at constraining the existing models (e.g. Edwards et al.
1994; Fernandez et al. 1995; Reipurth et al.
1996; Muzerolle et al. 1998b; Alencar & Basri
2000).
A wind scenario for the formation of the hydrogen emission lines
observed in T Tauri stars was much discussed over past
years. The predominance of blueshifted absorptions in the
H
line profiles, the success in modeling some of
these observations with wind models, together with the unequivocal
presence of winds in CTTS has been the main driving force behind
this interpretation. Despite the recent shift towards an
interpretation based on the magnetospheric accretion model, the
debate is far from settled, as it is clearly seen in Alencar &
Basri (2000). Also, for a given star, both winds and
accretion flows may contribute to hydrogen line emission, with
different lines being formed in distinct regions and/or affected
differently by those regions.
The successes and failures of the models in explaining hydrogen line
profiles in CTTS have been based on the study of lines from
the Balmer series. The exception is the work by Najita et al. (1996) where the Br
emission line is studied for
a very small number of objects.
The strongest hydrogen lines in the near infrared part of the
electromagnetic spectrum are lines from the Paschen and Brackett
series. These lines arise in higher energy levels when compared to
the lower Balmer lines (e.g. H
and
H
)
and they are generally optically thinner than
the latter. One expects them to form at different depths into a
cloud of hydrogen. Near infrared hydrogen lines should form in the
denser parts of the circumstellar envelope of CTTS and, in
particular, they should trace infalling material in a
magnetospheric accretion scenario. The development of near
infrared high spectral resolution high sensitivity spectroscopy
allows the study of profiles of near infrared hydrogen lines. Such
a study imposes strong constraints on models and contributes to
the understanding of the origin of hydrogen emission in
CTTS.
The structure of the paper is as follows: observations and data reduction are described in Sect. 2, the line spectra are presented in Sect. 3, classification of the line profiles is done in Sect. 4, the parameters that characterize the line profiles and their statistics are discussed in Sects. 5 and 6. The results are compared with other observations in Sect. 7, and with models in Sect. 8. The results are discussed in Sect. 9 and conclusions are presented in Sect. 10.
The observations were carried out at the United Kingdom
Infrared Telescope (UKIRT) during two observing runs. The first took
place between 2-5 October 1994 (UT), the echelle grating, long (300
mm) focal length camera and
pixel SBRC InSb detector of the
Cooled Grating Spectrometer 4 (CGS4) were
used. The second run was held between 15-17 December 1995 (UT) and CGS4 was
used with the echelle grating short (150 mm) focal length camera and
pixel SBRC InSb detector. Spectra were obtained centered
at Pa
(1.28215
m) and at Br
(2.16611
m) for a sample of T Tauri stars mainly from the
Taurus-Auriga complex. The observed stars are listed
in Table 1, with a "
'' indicating whether Pa
and/or Br
spectra was taken.
Star | RA(1950) | DEC(1950) | Type | Spec.Ty. |
![]() |
Pa
![]() |
Br
![]() |
V773 Tau | 4 11 7.29 | +28 04 41.2 | WTTS | K3 V | +16b | ![]() |
![]() |
FM Tau | 4 11 7.82 | +28 05 18.8 | CTTS | M0 | +16: (i) | ![]() |
![]() |
FN Tau | 4 11 8.61 | +28 20 26.9 | CTTS | M5 | +16; (ii) | ![]() |
|
CW Tau | 4 11 11.34 | +28 03 27.2 | CTTS | K3 V | +14.5a | ![]() |
![]() |
FP Tau | 4 11 43.50 | +26 38 57.5 | CTTS | M4 V | +22b | ![]() |
![]() |
CY Tau | 4 14 27.67 | +28 13 28.6 | CTTS | M1 V | +19.1a | ![]() |
|
DD Tau | 4 15 25.10 | +28 09 14.6 | CTTS | M1 | +28c | ![]() |
|
Hubble 4 | 4 15 40.89 | +28 12 54.0 | WTTS | K7 | +15.0a | ![]() |
![]() |
BP Tau | 4 16 8.61 | +28 59 15.3 | CTTS | K7 V | +15.8a | ![]() |
![]() |
LkCa 7 | 4 16 35.78 | +27 42 28.1 | WTTS | K7,M0V | +16.4a | ![]() |
![]() |
DE Tau | 4 18 49.84 | +27 48 05.2 | CTTS | M2: V | +14.9a | ![]() |
|
RY Tau | 4 18 50.85 | +28 19 35.0 | CTTS | K1IV,V | +16.4 | ![]() |
![]() |
FS Tau | 4 18 57.63 | +26 50 30.5 | CTTS | M1 | +16: (iii) | ![]() |
|
T Tau | 4 19 4.21 | +19 25 05.4 | CTTS | K0IV,V | +19.1a | ![]() |
![]() |
DF Tau | 4 23 59.63 | +25 35 41.7 | CTTS | M0,1 V | +12b | ![]() |
![]() |
DG Tau | 4 24 1.01 | +25 59 35.5 | CTTS | M ? | +16: (iv) | ![]() |
![]() |
DI Tau | 4 26 38.00 | +26 26 20.1 | WTTS | M0 V | +16.0a | ![]() |
|
IQ Tau | 4 26 47.67 | +26 00 16.3 | CTTS | M0.5 | +15.3a | ![]() |
|
FX Tau | 4 27 27.91 | +24 20 18.2 | CTTS | M1 | +16.8a | ![]() |
|
DK Tau | 4 27 40.48 | +25 54 59.0 | CTTS | K7 V | +15.3a | ![]() |
![]() |
ZZ Tau | 4 27 49.32 | +24 35 56.9 | CTTS | M3 | +15; (v) | ![]() |
|
HL Tau | 4 28 44.42 | +18 07 36.2 | CTTS | K7,M2? | +17: (vi) | ![]() |
![]() |
XZ Tau | 4 28 46.00 | +18 07 35.2 | CTTS | M3 | +17; (vii) | ![]() |
![]() |
HK Tau | 4 28 48.85 | +24 17 56.2 | CTTS | M0.5 | +16.3a | ![]() |
|
Haro 6-13 | 4 29 13.60 | +24 22 42.9 | CTTS | Cont. | +15; (v) | ![]() |
![]() |
GG Tau | 4 29 37.06 | +17 25 22.3 | CTTS | K7 V | +17.6a | ![]() |
![]() |
GH Tau | 4 30 4.79 | +24 03 18.3 | CTTS | M2,3 V | +18.4a | ![]() |
|
V807 Tau | 4 30 5.2 | +24 03 39 | CTTS | K7 V | +1 (viii) | ![]() |
![]() |
GI Tau | 4 30 32.33 | +24 15 3.1 | CTTS | K6 V | +18.1a | ![]() |
![]() |
GK Tau | 4 30 32.76 | +24 14 52.4 | CTTS | K7 V | +18.6a | ![]() |
![]() |
DL Tau | 4 30 36.02 | +25 14 24.0 | CTTS | K7 V | +16.0 (ix) | ![]() |
![]() |
AA Tau | 4 31 53.45 | +24 22 44.1 | CTTS | K7 V | +16.1a | ![]() |
![]() |
DN Tau | 4 32 25.68 | +24 08 52.3 | CTTS | M0 V | +16.1a | ![]() |
![]() |
HP Tau | 4 32 52.85 | +22 48 17.7 | CTTS | K3 | +17.7a | ![]() |
![]() |
DO Tau | 4 35 24.18 | +26 4 55.2 | CTTS | M0 V | +20:c | ![]() |
![]() |
VY Tau | 4 36 17.41 | +22 42 02.3 | var. | M0 V | +17.8a | ![]() |
|
DQ Tau | 4 43 59.99 | +16 54 40.1 | CTTS | M0,1 V | -4:c | ![]() |
|
Haro 6-37/c | 4 44 5.90 | +16 57 19.2 | CTTS | K7,M0 | +19.5b | ![]() |
![]() |
DR Tau | 4 44 13.20 | +16 53 23.8 | CTTS | Cont. | +16.7 (x) | ![]() |
![]() |
DS Tau | 4 44 39.07 | +29 19 56.2 | CTTS | K5 V | +16.3a | ![]() |
![]() |
UY Aur | 4 48 35.71 | +30 42 13.6 | CTTS | K7 V | +18b | ![]() |
![]() |
GM Aur | 4 51 59.76 | +30 17 14.7 | CTTS | K3 V | +15.0a | ![]() |
![]() |
SU Aur | 4 52 47.84 | +30 29 19.4 | CTTS | G2 III | +16.0 | ![]() |
![]() |
RW Aur A | 5 4 37.69 | +30 20 13.9 | CTTS | K1 | +14:a | ![]() |
![]() |
GW Ori | 5 26 20.78 | +11 49 52.8 | CTTS | G5 | +33.6a | ![]() |
![]() |
YY Ori | 5 32 20.77 | -5 59 52.7 | CTTS | K5 V | +12:c | ![]() |
|
TW Hya | 10 59 30.08 | -34 26 07.4 | CTTS | K7 V | +6:c | ![]() |
|
V1331 Cyg | 20 59 32.21 | +50 09 55.5 | CTTS | Cont. | -15: (xi) | ![]() |
![]() |
DI Cep | 22 54 8.18 | +58 23 59.5 | CTTS | G8 V | -10b | ![]() |
![]() |
BM And | 23 35 12.41 | +48 07 35.9 | CTTS | K5 V | -15.2 (xii) | ![]() |
![]() |
The spectral coverage is approximately
at
Pa
and approximately
for
spectra from
the UT94 run and approximately
at
Pa
and approximately
at
Br
for spectra from the UT95 run.
All observed spectra were sampled twice per resolution element (
for UT94 and
for UT95). The slit size
used was for the UT94 run:
with pixel
size
in the dispersion direction and
in the
spatial direction for both Pa
and Br
setups; and for the UT95 run:
and
the pixel size
in the dispersion direction and
and
in the spatial direction, respectively for the
Pa
and Br
wavelength regions.
In both runs flat
field frames were obtained for flat fielding and a number of standard
stars were observed for correction of the atmospheric transmission and
of the instrumental response. During
the UT95 run a number of late type main sequence stars were also
observed. For more details on the observational
procedures and observing logs refer to Folha (1998).
The data reduction followed the general procedure described by Puxley et al. (1992) for the reduction of spectra obtained with
CGS4. It was carried out using CGS4DR, Figaro and IDL routines written
specifically for this work. Spectral images were masked to avoid bad
pixels and vignetted areas of the detector, de-biased and
flat-fielded. Sky subtraction was achieved by subtracting sky frames from
object frames. The spectra were optimally extracted using Figaro's
implementation of the optimal extraction algorithm developed by Horne
(1986) and, when needed, de-rippled using the
deripple_spectrum task in CGS4DR. CVF fringing affected spectra
of the Br
window obtained on the UT95 run and the
fringing pattern was removed by filtering the spectra using standard
Figaro tasks. Wavelength calibration was achieved by using OH airglow
emission lines, telluric
absorption lines and, in some cases, photospheric absorption
lines. The uncertainty in the wavelength calibration is: 8-9
at Pa
and 4-9
at Br
spectra from UT94 and 6-8
at Pa
and 14-17
at Br
from UT95. Correction of the atmospheric
transmission and instrumental
response was achieved by dividing the spectra of the target stars by the
spectrum of one of the observed standard stars and multiplying the result
by a black body spectrum of the appropriate effective
temperature. Finally, spectra had their continuum normalized to unity by
dividing the spectra by a cubic spline fit to the continuum. For more
details on the data reduction procedure refer to Folha (1998).
A number of photospheric lines are seen in absorption in the observed
Pa
and Br
spectra of several T Tauri
stars (see Folha & Emerson 1999, hereafter FE99). Some of the narrow
photospheric absorption lines fall on top of the Pa
line,
changing its intrinsic profile. Correct analysis of the line profiles require
the photospheric component of the observed spectra to be removed.
This was achieved, after computing the veiling in the observed wavelength
ranges, by following the procedure described in Edwards et al.
(1994) to obtain "residual'' profiles (see FE99 for results on near
infrared veiling).
![]() |
Figure 1:
Pa
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Figure 1: continued |
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Figure 1: continued |
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Figure 1: continued |
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Figures 1 and 2 show the Pa
and Br
spectra of the observed T Tauri stars,
arranged by the type of the line profile, from type I to type IV (as
defined in Sect. 4 below) and for each type by
alphabetical order of the name of the star. For the stars for which
the veiling could be determined (FE99) the line profiles presented
here are "residual'' profiles, as defined by Edwards et al.
(1994). The small number of profiles obtained in UT94 are not
"residual'' profiles, as veiling measurements were not
possible with that data set.
Inspecting the line profiles one concludes that of the 49 stars observed
at Pa
,
42 show the line in emission and of the 36 stars
observed at Br
,
30 show the line in emission.
The stars without Pa
in emission are: DI Tau,
FN Tau,
Hubble 4, LkCa 7, V773 Tau, VY Tau
and ZZ Tau, of which DI Tau, Hubble 4, LkCa 7 and V773 Tau are
classified as Weak Line T Tauri stars (hereafter WTTS) in the Herbig & Bell
Catalogue (Herbig & Bell 1988).
The stars
without Br
in emission are: DN Tau,
FP Tau, GI Tau,
Hubble 4, LkCa 7
and V807 Tau of which Hubble 4
and LkCa 7 are WTTS.
From the stars observed both at Pa
and Br
some display emission in one line but not in the other.
Those
with emission
at Pa
but not at Br
are DN Tau,
FP Tau, GI Tau and V807 Tau. V773 Tau
displays emission at Br
but not at Pa
.
While the former is not surprising, the latter would be puzzling if
it were not for the highly variable characteristics of V773 Tau
(Feigelson et al. 1994).
The lines that appear in emission display a variety of profiles. Although no
two profiles are the same, many show similar properties and it is useful to
separate them according to the general shape of the line profile.
Type | % | No. Stars | Star name |
I | 53% | 20 | AA Tau BM And CY Tau |
DD Tau DE Tau DF Tau | |||
DG Tau DK Tau DL Tau | |||
DN Tau GG Tau Haro 6-13 | |||
Haro 6-37 HL Tau HK Tau | |||
T Tau UY Aur V1331 Cyg | |||
V807 Tau XZ Tau | |||
II B | 3% | 1 | CW Tau |
II R | 5% | 2 | DR Tau GW Ori |
III B | 0% | 0 | - |
III R | 5% | 2 | RY Tau SU Aur |
IV B | 0% | 0 | - |
IV R | 34% | 13 | BP Tau DO Tau DS Tau |
FM Tau FP Tau FS Tau | |||
GI Tau GK Tau GM Aur | |||
HP Tau IQ Tau RW Aur | |||
YY Ori |
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Figure 2:
Br
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Figure 2: continued |
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Figure 2: continued |
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The Reipurth et al. (1996) classification scheme divides the line profiles into four main types: type I profiles are generally symmetric showing no evidence for absorption features or only very slight influence from those; type II profiles show two peaks with the intensity of the second peak exceeding half the strength of the main peak; type III profiles show two peaks with the intensity of the second peak being less than half the strength of the main peak and finally type IV profiles, which show an absorption feature beyond which no emission is seen. To types II, III and IV the letters B or R are appended, depending on the location of the secondary peak/absorption feature relative to the main peak: if blueshifted a B is appended, if redshifted an R is added. Note that type IV B correspond to normal P Cygni profiles and type IV R correspond to inverse P Cygni profiles (henceforth IPC).
Type | % | No. Stars | Star name |
I | 72% | 18 | BP Tau DG Tau DI Cep |
DL Tau DO Tau DR Tau | |||
FM Tau GG Tau GM Aur | |||
GW Ori Haro 6-13 HL Tau | |||
T Tau TW Hya UY Aur | |||
V1331 Cyg V773 Tau XZ Tau | |||
II B | 0% | 0 | - |
II R | 8% | 2 | RY Tau SU Aur |
III B | 0% | 0 | - |
III R | 0% | 0 | - |
IV B | 0% | 0 | - |
II R | 20% | 5 | BM And CW Tau DF Tau |
HP Tau RW Aur |
The main conclusions that can be taken from Tables 2 and 3
are the following: most line profiles are
generally symmetric, especially the Br
line profiles,
where 72% of the profiles are classified as type I; both Pa
and Br
line profiles lack blueshifted absorptions (only one
star - CW Tau - displays any sort of absorption in the blue wing and
only in Pa
); nearly one third of the
Br
profiles have redshifted absorptions, with this
number being even higher for the Pa
profiles, of which 44%
display redshifted absorptions.
Star |
![]() |
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L/C |
(
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(
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(
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||
BP Tau |
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170 | 0.96 |
DO Tau |
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175 | 0.93 |
DS Tau |
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275 | 0.97 |
FM Tau |
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150 | 0.92 |
FS Tau |
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250 | 0.92 |
FP Tau |
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45 | 0.93 |
GI Tau |
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240 | 0.91 |
GK Tau |
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115 | 0.88 |
GM Aur |
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250 | 0.90 |
HP Tau |
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200 | 0.91 |
IQ Tau |
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200 | 0.97 |
RW Aur |
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300 | 0.95 |
YY Ori | ![]() |
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280 | 0.81 |
For a few stars the Pa
and/or Br
lines were observed more than once (refer to Sect. 9.4).
The classification
presented above refers to
data obtained during the second observing run, i.e. December 1995,
only. This is justified since, with the exception of the WTTS
LkCa 7, all stars observed during UT94 were also observed
during the 1995 campaign but with much higher
signal-to-noise.
The subject of variability in the line profiles will be left for
Sect. 9.4, however a few remarks are relevant to
the present discussion. The amount of variation present in the
Pa
line profile of the stars observed twice during
the 1995 campaign is such that the type of line profile did
not change. However, for some stars, significant changes are seen
in spectra taken 14 months apart, i.e. in October 1994 and
in December 1995. These changes do not significantly alter the
statistics presented, however they should be borne in mind.
Star |
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L/C |
(
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(
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(
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||
BM And |
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60 | 0.94 |
CW Tau |
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160 | 0.96 |
DF Tau |
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120 | 0.98 |
HP Tau |
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185 | 0.92 |
RW Aur |
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275 | 0.96 |
A method devised to measure these parameters and to estimate the associated uncertainties is described in Folha (1998).
The lines are very broad, with mean FWHM of
for Pa
and
for
Br
.
Line peaks are, generally, slightly blueshifted and very
rarely redshifted. The blueshift is, in most cases, very slight,
only seldom exceeding
for Pa
and
for Br
.
In fact, most lines
have their peaks located within
of the rest
velocity.
Typical parameters for each type of line profiles and comparison between them are discussed in the following sections.
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Figure 3:
Top panel: distribution of the Full Width at Half
Maximum for Pa
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Type I lines are broad, with FWHM ranging between 100 and
(Fig. 3), and slightly blueshifted
but seldom redshifted (Fig. 4).
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Figure 4:
Top panel: distribution of the line peak velocity for
Pa
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Figure 5:
Top panel: distribution of the asymmetry
factor, Af, for Pa
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The distributions of the maximum velocities observed in the blue and
red wings of type I profiles, are shown in Fig. 6 as
solid lines.
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Figure 6:
Distribution of the maximum velocities seen in the line
wings. Top panel: data for Pa
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Type IV R line profiles correspond to inverse P Cygni (IPC)
profiles and therefore are characterized by the presence of a
redshifted absorption feature. Typical velocities for these
are between
and
,
while their EWs can be as low as
for DS Tau's Pa
profile and as
large as
for YY Ori's
Pa
profile.
For an easier comparison between the observed properties of type IV R and of type I line profiles (which constitute the majority of the profiles), the distributions of the various observed parameters for IPC profiles are plotted with those corresponding to type I profiles in Figs. 4 to 6.
The observational characteristics of IPC profiles are similar
to those of type I profiles except for factors related to the
redshifted absorption, most notably apparent in the Af and
the maximum velocity observed for the red wing of the profile,
as can be clearly seen in Figs. 5 and
6. For IPC Pa
lines, the median Af is 1.67 (cf. median Af =1.16 for type I
Pa
profiles). For Br
IPC lines,
only one (in five) has a value for Af in the interval 0 to
5 (see Fig. 5). Of the remaining IPC
profiles, three have negative values for Af (due to the
larger equivalent width of the absorption than that of the
emission redward of the rest velocity) and one has an Af of
around 5.8. Also, the red wing in IPC profiles extends to less
than
for about 60% of the lines. In
contrast, type I profiles have red wings extending to such
low velocities for less than 35% of the lines.
The line widths of IPC and type I profiles, as measured by the FWHM,
are not significantly different. However, the former tend to be
narrower than the latter, with a larger percentage of stars of the
former type occupying the 150 to
bin in
Fig. 3. Also, the distributions of the line peak velocity
for IPC and type I profiles are similar, with a significant percentage
of profiles peaking just blueward of the system's rest velocity (see
Fig. 4). However, one should note that unlike the
distribution for type I Br
profiles, more than half
of the Br
IPC profiles have their line peaks
occurring between about
and
(Fig. 4). Finally, the distributions
of the maximum velocity observed at the blue wing of both IPC and
type I profiles are similar, with that for Br
IPC
profiles slightly shifted to higher velocities (Fig. 6).
Hamann et al. (1988) present high resolution, but relatively
low signal-to-noise, Br
spectra of DG Tau,
GW Ori, HL Tau, SU Aur and T Tau,
all of which are in the sample studied here. Despite the low
signal-to-noise ratio, it can be seen that the line profiles displayed
in their Fig. 1 have the same characteristics of the ones presented
here. Even SU Aur, seen with a slightly redshifted
absorption feature, is somewhat similar to the one shown in our Fig. 2.
Comparison with published line profiles other than the ones mentioned above is very difficult due to the much lower spectral resolutions and signal-to-noise ratios. Examples of those are the data shown in Giovanardi et al. (1991) and Evans et al. (1987) where the line profiles are defined by only a few spectral points.
Photometric standard stars or simultaneous photometry is not available for the data set presented here, hence reliable line fluxes could not be determined. Comparison with line fluxes available in the literature is, therefore, not possible. Variability in T Tauri stars, in particular in the NIR magnitudes, implies that any lines fluxes determined using non-simultaneous photometry in conjunction with the measured equivalent widths, would yield unreliable results.
Edwards et al. (1994) also present high resolution
observations of higher members of the Balmer series (from
H
to H
)
for a sample of 15 TTS.
These authors show that redshifted absorptions are seen frequently
in residual line profiles of the higher Balmer lines
(henceforth HBLs). Of the 8 TTS shown in Edwards et al.
(1994) to display IPC structure at least in one of the
Balmer lines presented there, 5 also display an IPC profile at
Pa
and/or at Br
.
GM Aur
displays IPC structure in Pa
but not in any of the
Balmer lines shown in Edwards et al. (1994). Furthermore,
for the Balmer lines, both redshifted and blueshifted absorptions
can coexist on the same line profile (e.g. DF Tau in
Edwards et al. 1994 and Alencar & Basri 2000).
That is not seen in any of the NIR line profiles presented here.
The Af of the HBLs falls mostly between 1 and 2.5.
Similar values for this parameter are found for Pa
and Br
.
The vast majority of the HBLs with
redshifted absorption features have them located between about
200 and
.
Only two of the objects in
the Edwards' sample have the centre of the absorption below
.
When compared to the HBLs, the NIR IPC
profiles tend to have the centre of the absorption feature located
at lower velocities, often below
(see
Tables 6 and 7). Furthermore,
the NIR lines analysed here tend to have line wings less extended
than those of the HBLs by about
.
Visual inspection of the Balmer profiles in Edwards et al.
(1994) reveal that most line wings extend up to
,
especially in the blue. From Fig. 6, we see that wings in NIR lines extend typically to
about
.
Also, IPC HBLs tend to have
line emission beyond the redshifted absorption feature. In
NIR lines, emission seems to stop blueward of the redshifted
absorption feature.
In summary, while the NIR lines show very different properties
from those of H
,
in some aspects they are similar
to higher Balmer lines. They are not as wide as the latter,
nor seem to be so much influenced by outflowing material. Like the
higher Balmer lines, the NIR lines are more prone to IPC
structure than H
.
The different velocities
at which the redshifted absorption feature occurs in IPC higher
Balmer lines and in IPC NIR lines should provide constraints on
models that hope to explain the formation of hydrogen lines in
TTS.
The above models concentrate on computing line profiles for Balmer
lines. Explicit model Pa
line profiles are not found
in the literature. Results for Br
lines are found in
Hartmann et al. (1990) and in Muzerolle et al. (1998a) only.
The model Br
profile shown in Hartmann et al. (1990) is the result from a non-isothermal wind with a spherically
symmetric steady flow (their model 12). Comparing the model line
profile with the high resolution observations presented here show no
resemblance between them. The model profile peaks at a redshifted
velocity larger than
while observed profiles
tend to peak at slightly blueshifted velocities. Also, the Af
of the model profile is clearly smaller than unity, again in
disagreement with the observations. Other near infrared lines
computed by Hartmann et al. (1990) (e.g. Pa
and
Br
)
show either a flat top with the line peak being
redshifted or to display a very prominent and broad blueshifted
absorption (P Cygni profile), again unlike any of the observed NIR
lines shown here.
Muzerolle et al. (1998a) show model Br
line
profiles, arising in a magnetospheric accretion scenario, for four
different viewing angles: inclinations of 10, 30, 60 and 75 degrees.
Full widths at half maximum range from
for
the lowest inclination to
for the highest
inclination. These are narrower than the typical observed lines,
which as we have seen in Sect. 6.2 above, display FWHM
mostly between 100 and
with about 60%
of the lines having FWHM between 150 and
.
Line wings in the model profiles extend to significantly
lower velocities than in the observed profiles, especially for the
lower inclination models. The 10 and 30 degrees inclination model
results show the blue wing extending up to
and the red wing extending to about
.
In
contrast, all of the observed blue wings and most of the observed red
wings are significantly more extended than those in the model profiles
(see distributions in Fig. 6). Furthermore, contrary to
the model predictions, the red wing in observed profiles is not more
extended than the blue wing. For the 60 and 75 degrees inclination
models (the model IPC profiles) the blue wing extends up to about
,
while the red wing extends up to
.
The redshifted absorption feature extends from
there up to about
.
These velocities match
reasonably well those observed for the maximum observed velocity of
the IPC profiles, but once again, fail at explaining the velocities
observed in the blue wing of the profiles. Another difference between
the model predicted and the observed profiles lies in the maximum line
intensity. Model profiles show normalized intensities of at least
about 1.8. Such high intensities are not observed even in TTS with
strong Br
emission. Quantitatively, judging from the
model Br
profiles presented in Muzerolle et al. (1998a), models do not match the observations very well.
However, a greater exploration of the model's parameter space is
required to assess this matter fully.
From a qualitative standpoint, Muzerolle et al. (1998a)'s
results for models with higher inclinations yield
Br
profiles which are IPC in shape and which
resemble some of the observed line profiles (e.g. DF Tau,
RW Aur), even if in quantitative terms they are
different (e.g. in their widths and peak intensities). The double
peaked profiles, characteristic of lower inclination models, are
not observed at Br
in the sample of stars
discussed here. None of the model Br
profiles
presented by Muzerolle et al. (1998a) is a type I profile.
These type of profiles constitute the vast majority of the
observed ones.
The presence of accretion/wind flows in TTS is signaled by
redshifted/blueshifted absorption features respectively dipping
below the continuum in their line profiles. Whenever Pa
and
Br
display absorption features they are
redshifted and dip below the continuum (CW Tau's
blueshifted absorption at Pa
is the exception).
These features must arise from infalling gas. The velocities
corresponding to those features are in good agreement with those
expected for free-falling material from a few stellar radii out.
In the context of the magnetospheric accretion scenario, these
velocities can be combined with an observational determination of
the stellar radius in order to obtain lower limits for the stellar
mass. Bonnell et al. (1998) used velocities provided by
Pa
and Br
data presented here to
obtain such mass estimates.
Modeling carried out in the context of magnetospheric
accretion (e.g. Hartmann et al. 1994) shows that
the presence of a redshifted absorption results from seeing
the infalling gas against the hot region where the accretion shock
occurs. The geometry of the system and the contrast between the
temperature of the infalling material and that of the shock region
are pivotal in determining whether an absorption appears in
the line profile and, if so, the strength of the absorption
feature. In fact, they seem to be much more important than the
value of the accretion rate itself. It is thus not
surprising that plotting accretion rates from Hartigan et al.
(1995) and from Gullbring et al. (1998) against the EW of the
absorption feature in the IPC Pa
line profiles
shows no significant correlation between the two quantities.
Also, no obvious differences in accretion rates are found between
stars displaying IPC or type I profiles. As suggested from the
modeling, the lack/presence of redshifted absorption in some
line profiles does not seem to have an obvious relation with stars
accreting matter at different rates or, at least, the accretion
rate is not the most important factor determining the actual shape
of the line profiles.
To investigate the role of the system's geometry in determining the
shape of the line profile we used inclinations available in the
literature (Bouvier et al. 1995) for the stars in our
sample. From the stars with inclinations in Bouvier et al. (1995), fifteen have Pa
IPC or type I
profiles and nine have Br
IPC or type I profiles.
Figure 7 shows the distribution of inclinations for
those stars in bins of
.
![]() |
Figure 7:
Distribution of inclinations for IPC and type I profiles. Top
panel - Pa
![]() ![]() |
Open with DEXTER |
![]() |
Figure 8:
Accretion
Rate vs. emission EW in Pa
![]() |
Open with DEXTER |
Since the presence of a redshifted absorption feature is also sensitive to the temperature of the accretion shock, one could, in principle, try to correlate the appearance of IPC structure in a given star with the temperature of its accretion shock. However, the latter is not observationally constrained well enough for individual objects.
What about the origin of line emission? According to Muzerolle et al. (1998a), if hydrogen line emission originates in infalling matter in the context of the magnetospheric accretion model, one should expect to see a correlation between emission line strength and accretion rate.
Plotting accretion/wind rates versus the EW of type I line profiles
show no apparent relation between the rate at which mass falls onto or
flows away from the stars and the strength of the Pa
or Br
(where the lines strenghts are normalised
to the continuum). Pearson's correlation coefficients are 0.44 and
0.41 respectively for
vs. EW and
vs.
EW. The number of data points in those plots (not shown) is 16 for
vs. EW and 14 for
vs. EW.
On the other hand, plotting accretion rates versus the EW of the
emission component in the
Pa
IPC lines (Fig. 8)
shows that there seems to be a trend in the sense
that stars with larger accretion rates tend to have larger EW in the
emission component of the IPC profiles.
![]() |
Figure 9:
Wind Mass Loss Rate vs. emission EW in Pa
![]() |
Open with DEXTER |
Plotting the wind mass loss rate, as determined by Hartigan et al. (1995), versus the EW of the emission component of the
IPC Pa
line profiles hints that no correlation exists
between the two quantities (see Fig. 9). This is
somewhat surprising given the positive correlation between mass
accretion rate and wind mass loss rate observed by Hartigan et al. (1995).
While the amount of emission in IPC profiles does seem to behave according to predictions from existing magnetospheric accretion models, the amount of emission in type I profiles does not seem to follow that trend. One possible explanation for the lack of correlation is that the main determinant of the strength of line emission in these NIR lines is not the amount of accreting matter. Alternatively, radiative transfer and/or geometrical effects may smooth out clear relationships between those two quantities.
Calvet & Hartmann (1992) were partially motivated by the fact that the Balmer line profiles observed for many T Tauri stars are generally symmetric and "centrally peaked'', i.e. type I profiles, which turned out to be very difficult to explain by wind models. That work and its development (Hartmann et al. 1994) show that infall models in a magnetospheric accretion scenario can produce the desired Balmer line profiles with no obvious signature for accretion. Hartmann et al. (1994) shows that profiles of hydrogen lines arising in magnetospheric accretion flows can be of type I but displaying a distinctive asymmetry due both to geometrical and radiative transfer effects. This asymmetry is in the sense that line profiles have Af's larger than unity and that line peaks occur at slightly blueshifted velocities. Edwards et al. (1994) use this asymmetry in Balmer lines as a signature for infall.
As discussed in Sect. 6.2 above, Pa
and Br
line profiles have their line peaks
ocurring at slightly blueshifted velocities and their Af's
are larger than unity (recall Figs. 4 and
5). Applying the same criteria used by Edwards et al.
(1994) one can argue that Pa
and
Br
type I profiles result from infall in a
magnetospheric accretion scenario. A different geometry and/or
different physical conditions in the accretion flow relative to
those considered by Hartmann, Calvet, Muzerolle and co-workers
seem to be necessary though. The observation of line profile
variability implying variable accretion (e.g. Johns & Basri
1995b) and the presence of localized accretion spots (e.g.
Unruh et al. 1998) indeed confirm that axially symmetric
accretion models, such as those considered by Hartmann and
co-workers, are not truly applicable in T Tauri stars. The virtue
of these models is that they do explain, in a qualitative fashion,
the general shape of the observed Pa
and
Br
line profiles. When quantitative comparisons
are done, the model Pa
and Br
lines are too narrow, by about
in both
FWHM and HWZI and also far too intense (by factors of a few in
peak intensity).
An alternative to infall models for the origin of line emission
are wind models. As thoroughly discussed by Calvet et al. (1992) and by Calvet & Hartmann (1992), the latter
models seem to have great difficulties in explaining type I
profiles. With the appropriate choice of parameters,
"stochastic'' wind models (Grinin & Mitskevich 1991 and
Mitskevich et al. 1993) are able to produce type I
profiles for lines of "intermediate optical depth''. These were
calculated with the intention of comparing the results with
observations of the infrared CaII triplet. Detailed calculations,
specific for the Pa
and Br
lines,
are needed in order to establish whether a clumpy structure for a
wind can produce type I profiles.
As referred to at the beginning of Sect. 3, a
number of narrow photospheric absorption lines fall on top of the
Pa
emission line (see FE99). These photospheric
lines are easily seen in many of the Pa
lines
observed. As they are photospheric in origin, the fact they
are seen imply that either Pa
is optically thin
or, alternatively, that the filling factor of the line emitting
region is only a small fraction of the stellar disk, allowing for
a direct view of the stellar photosphere.
How different is the information conveyed by Pa
with
respect to that given by Br
?
Almost half of the stars in the studied sample with line emission
either at Pa
or Br
have both
lines displaying similar characteristics.
The most conspicuous
differences
are stars that have IPC profiles at Pa
but type I at Br
or that display emission at
Pa
but no emission at all at Br
.
The former can result either from the physical conditions where
the lines are formed or from the non-simultaneity of the
Pa
and Br
line profiles (they
were obtained one day apart). Br
lines that have
type I profiles when the corresponding Pa
are IPC
tend to have Af's larger than the average for their class,
with more emission blueward of the line rest velocity. This tends
to point to the former explanation, with Br
optically thinner than Pa
and hence less prone to
display redshifted absorptions. If this is the case, it can
provide constraints on the physical conditions of the infalling
gas.
Lack of detected Br
emission while
Pa
is in emission shows that for some stars the
physical conditions in the hydrogen gas that surrounds them are
such that population of level 7 is not too significant while a
reasonable amount of atoms have level 5 populated.
A significant percentage of the Br
IPC lines tend
to peak more towards the blue than the Pa
IPC
lines (Sect. 6.2 above). The blueward shift is of
about
.
It is worthwhile noting that
amongst the Br
type I line profiles a
significant number also have the velocity of the line peak
shifted to the blue relative to
in the
Pa
lines (see Fig. 4). An explanation
for this shift would be that, due to their different optical
depths, Pa
and Br
sample regions
where gas moves at different speeds, resulting in line profiles
peaking at different velocities. Surprisingly the
magnetospheric accretion line profile modeling carried out by
Hartmann, Calvet, Muzerolle et al. does not show any
relative shift in the peak velocity between lines of very
different optical depths, such as H
,
H
,
H
,
and Br
.
A few TTS for which spectra are presented in this work were observed
more than once and some considerations regarding variations in the
observed NIR line profiles are due. The stars for which the NIR lines
were observed more than once are shown in Table 8 where
the dates of the observations are also indicated.
Observations in consecutive nights of the same line for the same star
were only carried out during the December 1995 run, only at
Pa
line and only for two nights. The night to
night variation observed is not very dramatic, in the sense that, for
these stars, the type of line profile did not change, despite changes
in equivalent width and/or in the position of absorption features.
For some of the stars, more drastic changes are observed to occur in
spectra taken roughly one year apart. DR Tau changes from
type I into type II R and GK Tau changes from type I into an
IPC line profile. The Br
lines of DR Tau
and GG Tau and the Pa
line of GI Tau
change significantly in terms of equivalent width (due to line profile
variations) but preserve the type of profile.
The type of variations that can occur in NIR lines are diverse
and must reflect the dynamical activity and/or changes in
the physical characteristics of the emitting region(s) (the change
of the Pa
profile of GK Tau from type I
to IPC surely indicates either that an accretion region came
into view or that accretion activity in the star changed between
the first and second observing runs). For a given star,
monitoring these lines over at least two rotation periods
(typically a week to a week and a half) would tell us how the
infalling matter behaves, since any variability displayed,
especially in a redshifted absorption component, should be
associated with an accretion flow. As an example, if a
magnetospheric accretion scenario is correct and the rotation axis
of the star is tilted relative to the magnetic axis one expects to
find a correlation between the line profiles and the phase of the
rotation period of the star (Johns & Basri 1995b). A
better understanding of how and where these lines are formed in
TTS certainly requires a variability study of these lines.
|
UT9410 | UT9512 | |||||
02 | 03 | 04 | 05 | 15 | 16 | 17 | |
DG Tau | - | Pa
![]() |
- | - | Pa
![]() |
Pa
![]() |
- |
DL Tau | - | Pa
![]() |
- | - | Pa
![]() |
Pa
![]() |
- |
DR Tau | - | Pa
![]() |
- | - | Pa
![]() |
Pa
![]() |
- |
- | Br
![]() |
- | - | - | - | Br
![]() |
|
GG Tau | - | - | - | Br
![]() |
- | - | Br
![]() |
GI Tau | Pa
![]() |
- | - | - | Pa
![]() |
- | - |
- | Br
![]() |
- | - | - | - | Br
![]() |
|
GK Tau | Pa
![]() |
- | - | - | Pa
![]() |
- | - |
- | Br
![]() |
- | - | - | - | Br
![]() |
|
RW Aur | - | - | - | - | Pa
![]() |
Pa
![]() |
- |
- | - | - | Br
![]() |
- | - | Br
![]() |
|
RY Tau | - | - | - | - | Pa
![]() |
Pa
![]() |
- |
SU Aur | - | - | - | - | Pa
![]() |
Pa
![]() |
- |
The Pa
and Br
lines are very wide
(FWHM
), slightly blueshifted and
with line wings extending to about
in the
blue and to about
in the red. The
Asymmetry Factor (Af) is slightly larger than one for most
cases, with the Br
distribution for this parameter
showing a larger spread than that of the Pa
distribution. The former also displays larger Afs than the
latter. A significant number of IPC Br
lines are
displaced to the blue by about
relative to
the IPC Pa
lines.
If the line emitting region for Pa
sits in front of
the whole stellar disk, the line is optically thin. Alternatively, the
filling factor of the line emitting region is small and the stellar
photosphere is also directly observed.
Comparing the data presented here and the results available in the
literature from models for the formation of the hydrogen lines
reveal that both wind and accretion models fail to explain the
observed line profiles in most cases. The accretion models,
computed in the context of the magnetospheric accretion scenario,
do provide a qualitative insight on how these lines might be
produced in the T Tauri stars' environment but fail under a
quantitative comparison. The models produce lines too narrow (by
FWHM), with wings extending to
velocities too small (by at least
)
and with far too high maximum normalized intensities (by factors
of a few). If nothing else, the discrepancies found between
observations and model line profiles hint that the axi-symmetric
models considered thus far are just a rough approximation to the
real accretion flows in T Tauri stars.
The redshifted absorption feature in the Pa
and
Br
IPC profiles must be formed in infalling material.
It is located at velocities of the order of the free-fall velocity
from a few radii out for a typical T Tauri star. There seems to be a
trend associating the amount of emission seen in the IPC profiles and
accretion rates, suggesting that lines with this type of profile
originate mostly, if not completely, from infalling material. A
similar trend does not seem to be present for type I profiles.
However, with the exception of the lack of redshifted absorption
feature, the latter display similar characteristics to the IPC
profiles. In particular, they are centrally peaked and slightly
blueshifted, characteristics that are very difficult to obtain in wind
models (Calvet & Hartmann 1992) but arise naturally in inflow
models due to absorption of infalling redshifted material.
Continuous wind models tend to produce lines with normal P Cygni
profiles but Pa
or Br
calculations
are seldom found in the literature. Stochastic wind models might be
able to produce profiles similar to type I but there are no specific
predictions for the Pa
and Br
lines.
The way in which winds affect the shape of these NIR lines should be
investigated further.
The data set presented here demonstrate that current knowledge about the formation of hydrogen lines in T Tauri stars is far from providing a detailed explanation for their characteristics and origin. Also, it provides a solid database with which model results can be compared. Hydrogen lines constitute one of the most important diagnostics available for the study of T Tauri stars and understanding their origin is of crucial importance. From a theoretical point of view, models have to simultaneously explain the near infrared lines and the Balmer lines, which as we have seen above convey different information. An observational effort to try to understand how and why the lines vary is also very important. Otherwise, we will only be trying to understand T Tauri stars and what gives rise to hydrogen emission lines from a single snapshot of what is, in reality, an ever changing system.
Acknowledgements
We thank the referee, Dr. Suzan Edwards, for insightful comments that significantly improved this paper. D.F.M. Folha acknowledges financial support from the "Subprograma Ciência e Tecnologia doQuadro Comunitário de Apoio''. This research has made use of the Simbad database, operated at CDS, Strasbourg, France. The United Kingdom Infrared Telescope, is operated by Joint Astronomy Centre on behalf of the U.K. Particle Physics and Astronomy Research Council.