TEXES Observations of Pure Rotational H2 Emission From AB Aurigae
arXiv:0704.1481v1  [astro-ph]  11 Apr 2007
Draft version November 20, 2021
Preprint typeset using LATEX style emulateapj v. 2/19/04
TEXES OBSERVATIONS OF PURE ROTATIONAL H2 EMISSION FROM AB AURIGAE
Martin A. Bitner
, Matthew J. Richter
, John H. Lacy
, Thomas K. Greathouse
, Daniel T. Jaffe
Geoffrey A. Blake
Draft version November 20, 2021
ABSTRACT
We present observations of pure rotational molecular hydrogen emission from the Herbig Ae star, AB
Aurigae. Our observations were made using the Texas Echelon Cross Echelle Spectrograph (TEXES)
at the NASA Infrared Telescope Facility and the Gemini North Observatory. We searched for H2
emission in the S(1), S(2), and S(4) lines at high spectral resolution and detected all three. By
fitting a simple model for the emission in the three transitions, we derive T = 670 ± 40 K and M
= 0.52±0.15 M⊕ for the emitting gas. Based on the 8.5 km s
−1 FWHM of the S(2) line, assuming the
emission comes from the circumstellar disk, and with an inclination estimate of the AB Aur system
taken from the literature, we place the location for the emission near 18 AU. Comparison of our
derived temperature to a disk structure model suggests that UV and X-ray heating are important in
heating the disk atmosphere.
Subject headings: circumstellar matter, infrared:stars, stars:planetary systems, protoplanetary disks,
stars:individual(AB Aur), stars:pre-main sequence
1. INTRODUCTION
Disks around young stars are a natural part of the
star formation process and, as the likely site of planet
formation, have generated much interest. A consider-
able amount has been learned about the structure of
circumstellar disks through the study of dust emission
and detailed modeling of the spectral energy distribu-
tions (SED) of these stars. However, if we assume the
standard gas-to-dust ratio for the interstellar medium,
dust makes up only 1% of the mass of the disk. Thus
there is an interest in developing direct tracers of the gas
component. Studies of gas in disks have so far focused
mainly on gas at either large radii (> 50 AU), using ob-
servations at submillimeter wavelengths (Semenov et al.
2005), or small radii, using observations of near-infrared
CO lines (Blake & Boogert 2004; Najita et al. 2003).
Mid-infrared spectral diagnostics may probe disks at in-
termediate radii (1-10 AU) (Najita et al. 2006).
Molecular hydrogen diagnostics are promising because
they trace the dominant constituent of the disk and so do
not rely on conversion factors to determine the mass of
the emitting gas. H2 has been observed in circumstellar
environments at ultraviolet (Johns-Krull et al. 2000) and
near-infrared (Bary et al. 2003) wavelengths. These ob-
servations trace hot circumstellar gas, or gas excited by
fluorescent processes, and are therefore difficult to trans-
1 Department of Astronomy, University of Texas at
Austin, Austin, TX 78712; mbitner@astro.as.utexas.edu,
lacy@astro.as.utexas.edu, dtj@astro.as.utexas.edu
2 Visiting Astronomer at the Infrared Telescope Facility,
which is operated by the University of Hawaii under Cooperative
Agreement no. NCC 5-538 with the National Aeronautics and
Space Administration, Science Mission Directorate, Planetary
Astronomy Program.
3 Physics Department, University of California at Davis, Davis,
CA 95616; richter@physics.davis.edu
4 Lunar and Planetary Institute, Houston, TX 77058;
greathouse@lpi.usra.edu
5 Division of Geological & Planetary Sciences, California
Institute of Technology, MS 150-21, Pasadena, CA 91125;
gab@gps.caltech.edu
late into gas masses. The pure rotational mid-infrared
H2 lines are useful probes because the level populations
should be in LTE at the local gas temperature, and so
line ratios allow determination of the excitation temper-
ature and mass of the warm gas. The line ratios of the
three low-J lines accessible from the ground are sensitive
to gas at temperatures of 200-800 K. Other lines in the
mid-infrared such as [NeII] at 12.8 µm (Glassgold et al.
2007), [SI] at 25.2 µm and [FeII] at 26 µm (Pascucci et
al. 2006) should be useful probes of gas in disks but their
interpretation requires detailed modeling. Molecular hy-
drogen emission does not trace the cold, optically thick
regions of disks, and therefore estimates of the amount
of warm gas based on H2 emission represent a lower limit
to the total disk mass. In order to see H2 emission aris-
ing from a disk, the warm gas must be either physically
separated from optically thick dust or at a different tem-
perature. These conditions can be met in the disk sur-
face layer or in gaps or holes in the disk. In either case,
gas can be physically separated (due to dust settling or
clearing) as well as thermally decoupled (due to lower
densities) from the dust and heated to a higher temper-
ature by X-ray and UV radiation (Glassgold et al. 2004;
Gorti & Hollenbach 2004; Nomura & Millar 2005).
From the ground, three H2 mid-infrared rotational
lines are accessible: S(1) (λ = 17.035 µm), S(2) (λ =
12.279 µm), and S(4) (λ = 8.025 µm). The high spec-
tral resolution achievable with the Texas Echelon Cross
Echelle Spectrograph (TEXES; Lacy et al. 2002) is cru-
cial for these observations in order to separate the lines
from nearby telluric features and to maximize the line
contrast against the dust continuum. Observations at
high spectral resolution also have the advantage that,
when coupled with information about the disk inclina-
tion, they allow an estimate for the location of the emit-
ting gas. For emission lines with small equivalent width,
ground-based observations with TEXES can be more
sensitive than Spitzer IRS observations.
AB Aur is one of the brightest and thus one of the
most well-studied of the Herbig Ae stars. It is located
http://arxiv.org/abs/0704.1481v1
2 BITNER ET AL.
at a distance of 144 pc based on Hipparcos measure-
ments (van den Ancker et al. 1998), has a spectral type
of A0-A1 (Hernández et al. 2004) and is surrounded by a
disk/envelope structure that extends to at least r ∼ 450
AU (Mannings & Sargent 1997). Observations at 11.7
and 18.7 µm by Chen & Jura (2003) show that AB Aur
varies in the mid-infrared. The mass of the star is 2.4
M⊙ with an age of 2-4 Myr (van den Ancker et al. 1998).
AB Aur has a SED which is well-fit by the passive irradi-
ated disk with puffed-up inner rim model of Dullemond
et al. (2001). There have been a wide range of inclina-
tion estimates for AB Aur (see Brittain et al. 2003 for an
extended discussion of this issue). Recent results by Se-
menov et al. (2005) derive an inclination of 17+6
−3 deg for
the AB Aur system by modeling millimeter observations.
We will assume this value for analysis of our data.
Several attempts have been made to observe molecular
hydrogen emission from the disk of AB Aur. Thi et al.
(2001) claimed the detection of H2 S(0) and S(1) emis-
sion associated with the AB Aur system using the In-
frared Space Observatory (ISO), but subsequent obser-
vations from the ground with improved spatial resolu-
tion were unable to confirm the emission (Richter et al.
2002). Richter et al. (2002) and Sheret et al. (2003) saw
some evidence of emission in the S(2) line but the data
quality was such that a definitive statement could not be
made. They saw no evidence of the S(1) line. In this
paper, we present definite detections of the S(1), S(2),
and S(4) lines.
2. OBSERVATIONS AND DATA REDUCTION
All of the data were acquired using TEXES in its high-
resolution mode on the NASA Infrared Telescope Facil-
ity (IRTF) between December 2002 and October 2004
and on Gemini North in November 2006 under program
ID GN-2006B-Q-42. Table 1 lists details of the obser-
vations. We nodded the source along the slit to remove
background sky emission. We observed flux and telluric
standards at each setting. Other calibration files taken
include blank sky and an ambient temperature black-
body used to calibrate the wavelength and to flatfield
the data.
The data were reduced using the standard TEXES
pipeline (Lacy et al. 2002) which produces wavelength-
calibrated one dimensional spectra. The TEXES pipeline
data reduction gives a first order flux calibration. How-
ever, because of the different slit illumination of the am-
bient temperature blackbody and point sources, we ob-
served flux standards to determine a correction factor.
We assumed that the slit illumination was the same for
the target as for the flux standard. Guiding and seeing
variations can introduce uncertainty into our flux cali-
bration.
3. RESULTS
Figure 1 shows flux-calibrated spectra for all three set-
tings taken at the Gemini North Observatory. Overplot-
ted on each spectrum is a Gaussian with integrated line
flux equal to that derived from the best-fit model as-
suming the emission arises from an isothermal mass of
optically thin H2 gas. The three lines were fit simulta-
neously while varying the temperature and mass of the
emitting gas in the model. For the simultaneous fitting,
we adopted the FWHM and centroid of the S(2) line for
all three lines. The quoted errors for the temperature
and mass are 1-σ based on the contour plot of the χ2
values. Our best-fit model based on Gemini data has
T = 670 ± 40 K and M = 0.52 ± 0.15 M⊕. A similar
analysis of our IRTF data, which consisted of a detec-
tion at S(2) and upper limits for S(1) and S(4), gives
T = 630± 60 K and M = 0.7± 0.2 M⊕. Though within
2-σ uncertainties, the lower measured S(2) flux observed
at Gemini than from the IRTF combined with the fact
that our slit is smaller on the sky at Gemini suggests we
may be resolving out some of the flux we see at the IRTF.
To investigate this possibility, we examined spatial plots
of our 2-D echellograms but found no clear evidence for
the existence of spatially resolved line emission.
Table 2 contains a summary of our results. We fit each
line individually with a Gaussian to determine centroid,
FWHM, and flux values. The lines are all centered near
the systemic velocity of AB Aur and have FWHM ∼10
km s−1. Roberge et al. (2001) find A
= 0.25 for AB
Aur, so we assume no extinction at our wavelengths. We
quote equivalent widths in addition to line fluxes due to
uncertainty in the determination of the continuum level.
In an attempt to test how sensitive our temperature and
mass estimates are to errors in the flux calibration, we fit
our data after normalizing to the continuum flux values
from ISO SWS observations of AB Aur (Meeus et al.
2001). The resulting temperature and mass do not differ
significantly from the values based on our internal flux
calibration. We derive T = 630 K and M = 0.71 M⊕
when normalizing to the ISO continuum values of 13.1
Jy at 8 µm, 21.0 Jy at 12 µm, and 30.8 Jy at 17 µm
before fitting. The upper limits for S(1) and S(4) from
the IRTF quoted in Table 2 are based on a Gaussian fit
at the expected position of each line plus a 1-σ error.
The 1-σ line flux errors were computed by summing over
the number of pixels corresponding to the FWHM of the
lines in regions of the spectrum with similar atmospheric
transmission.
Figure 2 shows a population diagram based on our
observations. The points marked are based on the line
fluxes derived from Gemini observations listed in Table 2,
and the error bars are 1-σ. The overplotted solid line is
not a fit to the population diagram; rather it is based on
the best-fit model temperature and mass. Even though
we only have three points, the deviation from a single
temperature (a straight line on this plot) is significant
and may indicate that the emission is coming from a mix
of temperatures, as expected for a disk with a radial tem-
perature gradient. We note that curvature in the popu-
lation diagram can also occur if the ortho-to-para (OTP)
ratio is different from our assumed value of 3. Fuente et
al. (1999) derived a an OTP ratio between 1.5-2 for H2
in a photodissociation region with gas temperatures of
300-700 K. However, the sense of the curvature in our
population diagram would require an OTP ratio greater
than 3 to explain the increased S(1) flux, which seems
unlikely. We discuss the derived temperature of the gas
and its location in the next section.
4. DISCUSSION
We have detected all three mid-IR H2 lines with good
signal to noise. The results are consistent with the pre-
vious upper limits and the 2-σ detection of the S(2) line
(Richter et al. 2002; Sheret et al. 2003). Our observations
OBSERVATIONS OF H2 EMISSION FROM AB AURIGAE 3
of all three lines allow us to put tighter constraints on the
temperature and mass of the emitting gas. The analysis
of our data has three components. First, we determine a
characteristic temperature of the emitting gas by fitting a
simple model to the three lines. Second, we locate the po-
sition in the disk where the emission arises by comparing
our observed line profiles to line profiles calculated from
a Keplerian disk model. Finally, in order to comment
on the structure of the disk, we compare our results to
the predictions of a well-established disk structure model
which assumes well-mixed gas and dust at the same tem-
perature and find that, under such assumptions, we do
not reproduce our observations. This argues for the need
for additional gas heating mechanisms which may come
in the form of X-ray and UV heating. The importance of
UV and X-ray heating in the surface layers of disks has
been demonstrated observationally by Bary et al. (2003)
and Qi et al. (2006).
The procedure used to derive a temperature of T =
670 ± 40 K and mass of M = 0.52 ± 0.15 M⊕ from a
simulaneous fit to the H2 S(1), S(2), and S(4) lines as-
sumes the gas is optically thin and in LTE at a single
temperature. If the emission arises in the atmosphere of
a disk where radial and vertical temperature gradients
exist, the derived temperature should be considered a
characteristic or average temperature. In fact, Figure 2
shows evidence that the emission arises in gas at a mix
of temperatures. It should be noted that radial tem-
perature gradients in a disk also affect the detectability
of various lines. Since we are most sensitive to narrow
emission lines, the S(1) line, originating in cooler gas
farther out in the disk, is easier to detect than emission
in the S(4) line. Emission arising very close to the star
from the inner rim of the disk would be more difficult for
us to detect at our high spectral resolution due to line
broadening.
Our derived temperature, 670 K, is in agreement with
the temperature found by Richter et al. (2002) who de-
rived T > 380 K based on S(2) and the upper limit on
S(1). Brittain et al. (2003) made observations toward AB
Aur at 4.7 µm of 12CO v = 1-0 emission. The low-J lines
in the resulting population diagram have a steeper slope,
interpreted as emission coming from cooler gas than the
high-J lines. Fits to the population diagrams assuming
LTE gas give T = 70 K for the cool gas and T = 1540 K
for the hot gas. The authors explain the hot CO emission
as coming from the inner rim of the disk, while the cool
emission originates in the outer flared part of the disk
after it emerges from the inner rim’s shadow. Blake &
Boogert (2004) also observed 4.7 µm CO emission from
AB Aur. Isothermal fits with T ∼ 800 K are able to
reproduce their line fluxes. The line widths for the CO
lines are generally broader than what we observed, sug-
gesting that the CO emission comes from smaller radii
or outflowing gas (Blake & Boogert 2004).
Our derived mass, 0.52 M⊕, is four orders of magni-
tude less than the total disk mass of AB Aur (Semenov
et al. 2005). This is partly due to the fact that AB
Aur’s disk is optically thick at mid-infrared wavelengths
(Sheret et al. 2003). Most of the gas in the disk is hid-
den by optically thick dust and is also too cool to emit
at our wavelengths. The mass we derive should therefore
be interpreted as a lower limit to the amount of warm
gas on the surface of the disk facing us.
By comparing the FWHM of our observed S(2) line
profile with computed line profiles from a simple Keple-
rian disk model, we derive an estimate of where the line
emission originates. We generate line profiles under the
assumption of emission from a Keplerian disk inclined at
17◦ (Semenov et al. 2005) orbiting a 2.4 M⊙ star (van
den Ancker et al. 1998) with equal emissivity at all points
within an annulus 2 AU in extent. We convolve the line
profiles with our instrumental broadening function plus
a thermal broadening function based on our derived tem-
perature, T = 670 K, and compare the resulting FWHM
of the line profiles to our observed S(2) line. Under these
assumptions, the line profiles imply that our emission
arises near 18 AU in the disk.
We used parameters for the disk based on the fit to
AB Aurigae’s SED from Dullemond et al. (2001) as input
to the vertical disk structure model of Dullemond et al.
(2002) in order to derive a grid of temperature and den-
sity values as a function of disk radius and height above
the midplane. The Dullemond et al. (2001) model does
an excellent job fitting the SED, although we note that
their derived inclination angle, 65◦, is significantly larger
than our assumed value of 17◦. As noted in Dullemond
et al. (2001), the large discrepancy between their derived
inclination and what most observations suggest is proba-
bly due to the inclusion in their SED model of a perfectly
vertical inner rim which may not be physically accurate.
We input this radial and vertical temperature and den-
sity profile into an LTE radiative transfer program. We
assume the gas is in LTE with the dust and compute the
emergent spectrum. The emission lines produced by our
model under these conditions are much weaker than our
observations. This is not especially surprising given our
assumption of well-mixed gas and dust at the same tem-
perature and suggests the need for additional gas heating
mechanisms.
The vertical structure model of Dullemond et al. (2002)
is mainly concerned with modeling the dust emission
from disks. By assuming equal gas and dust temper-
ature, the gas only reaches a temperature of ∼ 260 K
at the top of the disk at 18 AU, much cooler than our
derived temperature of 670 K. To explain the observed
emission, an additional mechanism is needed to heat the
gas. X-ray and UV heating are likely possibilities which
can heat the gas to temperatures significantly hotter than
the dust. Glassgold et al. (2007) computed the gas tem-
perature in X-ray irradiated disks around T Tauri stars.
At 20 AU, the temperature can reach over 3000 K at the
top of the disk before dropping to 500-2000 K in a transi-
tion zone and then to much cooler temperatures deep in
the disk. Nomura & Millar (2005) considered UV heat-
ing of protoplanetary disks and the resulting molecular
hydrogen emission. The temperature in their model at
10 AU reaches over 1000 K and predicts an H2 S(2) line
flux of 2×10−15 ergs s−1 cm−2, somewhat lower than our
observed flux, but their model was for a less massive star
and disk. H2 emission could also arise in regions where
there is spatial separation of the gas from the dust due
to dust settling or coagulation of dust into larger parti-
cles. The SED of AB Aur shows no evidence for gaps in
the AB Aur disk and we are unable to make definitive
statements based on our observations.
A UV-heated H2 layer in a protostellar disk ought to
exhibit at least some similarities to the comparable hot
4 BITNER ET AL.
layer at the surface of photodissociation regions (PDRs)
with high densities and strong radiation fields. We there-
fore compare our observations of AB Aur with earlier
mid-IR H2 spectroscopy of the Orion Bar PDR. The UV
field expressed in terms of the mean interstellar radiation
field that was adopted by Allers et al. (2005) (3× 104) is
somewhat smaller than that expected at 18 AU around
an A0 star (∼ 105). Jonkheid et al. (2007) have com-
puted the strength of the UV field throughout a disk
around a Herbig Ae star finding values between 105 - 106
near the disk surface at 20 AU. Our derived tempera-
ture, 670 K, is in the range of temperatures (400-700 K)
derived for the Orion bar PDR based on the H2 S(1),
S(2), and S(4) lines (Allers et al. 2005). As shown in
Allers et al. (2005), the line ratios are determined by a
complex mix of the temperature gradient, H2 abundance
and dust properties. As another check of the plausibility
that the surface of AB Aur’s disk is similar to a PDR,
we computed the surface brightnesses of our three lines.
Assuming our observed emission arises from an annulus
2 AU in extent centered at 18 AU, we derive surface
brightnesses of ∼ 10−5 ergs cm−2 s−1 sr−1, roughly an
order of magnitude lower than those seen by (Allers et
al. 2005). However, Allers et al. (2005) estimated that
their line intensities were brightened by a factor of 10
since the Orion Bar PDR is seen nearly edge-on making
our surface brightnesses consistent with each other.
We thank the Gemini staff, and John White in par-
ticular, for their support in getting TEXES to work on
Gemini North. We thank Rob Robinson for useful dis-
cussions on the analysis of our data. The development of
TEXES was supported by grants from the NSF and the
NASA/USRA SOFIA project. Modification of TEXES
for use on Gemini was supported by Gemini Observa-
tory. Observations with TEXES were supported by NSF
grant AST-0607312. This work is based on observa-
tions obtained at the Gemini Observatory, which is op-
erated by the Association of Universities for Research
in Astronomy, Inc., under a cooperative agreement with
the NSF on behalf of the Gemini partnership: the Na-
tional Science Foundation (United States), the Particle
Physics and Astronomy Research Council (United King-
dom), the National Research Council (Canada), CON-
ICYT (Chile), the Australian Research Council (Aus-
tralia), CNPq (Brazil) and CONICET (Argentina).
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OBSERVATIONS OF H2 EMISSION FROM AB AURIGAE 5
TABLE 1
Observing Parameters
Date Telescope Linea Resolving Powerb Slit Widthc Integration Time
(R ≡ λ/δλ) (arcsec, AU) (s)
Dec 2002 IRTF S(1) 60,000 2.0, 288 2591
Dec 2002 IRTF S(2) 85,000 1.4, 202 4792
Dec 2003 IRTF S(2) 81,000 1.4, 202 2072
Oct 2004 IRTF S(4) 88,000 1.4, 202 8549
Nov 2006 Gemini S(1) 80,000 0.81, 117 2007
Nov 2006 Gemini S(2) 100,000 0.54, 78 2072
Nov 2006 Gemini S(4) 100,000 0.54, 78 1295
aS(1) is at 587.0324 cm−1. S(2) is at 814.4246 cm−1. S(4) is at 1246.098 cm−1.
bThe resolving power is deduced from observations of stratospheric emission from Titan
or Saturn at similar frequencies or unresolved emission lines in a low pressure gas cell.
cThe slit width in AU is calculated at AB Aur’s distance of 144 pc.
6 BITNER ET AL.
TABLE 2
Summary of Results
Telescope λ Fν Line Flux Equivalent Width FWHM
(µm) (Jy) (10−14 ergs s−1 cm−2) (km s−1) (km s−1)
8 12.9 (1.1) < 1.2 ... ...
IRTF 12 13.3 (0.4) 0.93 (0.25) 0.86 (0.23) 7.0
17 25.2 (1.8) < 1.1 ... ...
8 12.7 (0.4) 1.47 (0.34) 0.93 (0.21) 10.4
Gemini 12 14.7 (0.3) 0.53 (0.07) 0.44 (0.06) 8.5
17 24.6 (0.4) 1.1 (0.30) 0.76 (0.21) 9.0
OBSERVATIONS OF H2 EMISSION FROM AB AURIGAE 7
Fig. 1.— Regions of the three H2 pure rotational lines available from the ground. The data shown were obtained at Gemini North in
November 2006 under Gemini program GN-2006B-Q-42. Each spectrum has been flux-calibrated. The lines were individually fit with a
Gaussian to determine their centroid, flux, and FWHM. The lines are all centered near the systemic velocity of AB Aur and have FWHM
∼10 km s−1. Line fluxes are listed in Table 2. The overplotted Gaussians are based on the simultaneous fit and are not the individual fits
to each line. The three spectral regions were simultaneously fit with a model assuming the observed emission arises from a mass of H2 gas
at a single temperature. The best-fit model with T = 670 K and H2 mass = 0.52 M⊕ is overplotted.
8 BITNER ET AL.
Fig. 2.— Population diagram based on Gemini North observations. The points marked are based on fluxes derived by fitting a Gaussian
to each of the lines and shown with 1-σ error bars. The overplotted solid line is not a fit to the population diagram; rather it is based on the
best-fit model temperature and mass found by simultaneously fitting all three lines. The deviation of the points from the single-temperature
line is significant and suggests that the emission may arise from gas at a range of temperatures.
