# Near-IR H2 Emission of Protostars: Probing Circumstellar Environments111The
data presented herein were obtained at the W.M. Keck Observatory from
telescope time allocated to the National Aeronautics and Space Administration
through the agency’s scientific partnership with the California Institute of
Technology and the University of California. The Observatory was made possible
by the generous financial support of the W.M. Keck Foundation.

Thomas P. Greene  NASA Ames Research Center, M.S. 245-6, Moffett Field, CA
94035<EMAIL_ADDRESS>Mary Barsony22affiliation: Space Science Institute,
4750 Walnut Street, Suite 205, Boulder, CO 80301 Department of Physics and
Astronomy, San Francisco State University, 1600 Holloway Drive, San Francisco,
CA 94132<EMAIL_ADDRESS>David A. Weintraub  Department of Physics
and Astronomy, Vanderbilt University, Nashville, TN 37235
<EMAIL_ADDRESS>
###### Abstract

We present new observations of near-infrared molecular hydrogen (H2) line
emission in a sample of 18 Class I and flat-spectrum low mass protostars,
primarily in the Tau-Aur and $\rho$ Oph dark clouds. The line emission is
extended by up to several arcseconds (several hundred AU) for most objects,
and there is little night-to-night variation in line strength coincident with
the continuum point source. Flux ratios of H2 $v=2-1$ $S(1)$ and $v=1-0$
$S(1)$ lines are consistent with this emission arising in jets or winds in
many objects. However, most objects have only small offsets (under 10 km s-1)
between their H2 and photospheric radial velocities. No objects have line
ratios which are clearly caused solely by UV excitation, but the H2 emission
of several objects may be caused by UV or X-ray excitation in the presence of
circumstellar dust. There are several objects in the sample whose observed
velocities and line fluxes suggest quiescent, non-mechanical origins for their
molecular hydrogen emissions. Overall we find the H2 emission properties of
these protostars to be similar to the T Tauri stars studied in previous
surveys.

ISM: jets and outflows — stars: pre-main-sequence, formation — infrared: stars
— techniques: spectroscopic

††slugcomment: Accepted by ApJ on 8 October 2010

##  1 Introduction

Embedded low-mass protostars have been identified from their infrared (IR)
energy distributions for over two decades, but their high extinctions and
relatively small sizes make it very difficult to observe how radiation from
the central protostars interacts with their inner, pre-planetary circumstellar
disks. Recent high sensitivity, high resolution spectroscopic surveys have
revealed the detailed stellar properties of significant numbers of these
objects (Doppmann et al., 2005; White & Hillenbrand, 2004; Greene & Lada,
2002), but less is known about the physical conditions in their disks,
especially of the gaseous component, or about their inner disks where planets
might form and migrate. Here, we explore the gaseous component of the
innermost regions of self-embedded (Class I and flat-spectrum; hereafter FS)
protostars via high-resolution near-infrared (NIR) spectroscopy of the most
abundant molecule in proto-planetary disks, molecular hydrogen (H2). Near-
infrared H2 transitions are used, since the high extinctions to these sources
preclude UV spectroscopy.

The NIR ro-vibrational lines of H2 are good tracers of physical conditions in
inner circumstellar disks and winds close to protostars. There are several
different mechanisms that might be responsible for the production of NIR H2
emission lines in late-stage protostars. These include: i) shock heating of
the ambient medium by winds or jets, ii) X-ray heating, or iii) UV-heating,
level pumping and fluorescence.

Protostars must accrete mass at mean rates of $\sim 10^{-6}-10^{-5}M_{\sun}$
yr-1 to assemble themselves on time scales of several $10^{5}$ yr, and these
high rates produce significant UV flux when the accreting matter impacts the
stellar surface (e.g., Gullbring et al., 2000). Protostars and T Tauri stars
are also known to be strong and variable sources of X-ray emission (e.g.,
Imanishi et al., 2001, 2003; Güdel et al., 2007; Flaccomio et al., 2009). Both
UV and X-rays can excite vibrational states of H2, producing NIR line emission
(e.g., Gredel & Dalgarno, 1995; Nomura et al., 2007). Many protostars and T
Tauri stars also shed mass in jets that drive molecular outflows, and these
jets also frequently excite NIR vibrational H2 line emission (Zinnecker et
al., 1998).

The most commonly encountered excitation mechanism for NIR H2 emission
associated with protostars to date has been shock-excitation222See
http://www.jach.hawaii.edu/UKIRT/MHCat/ for an up-to-date listing of
“Molecular Hydrogen Emission-Line Objects (MHOs) in Outflows from Young
Stars.” (e.g., Davis et al., 2010), consistent with many self-embedded
protostars being associated with large-scale molecular outflows (e.g.,
Moriarty-Schieven et al., 1992). These outflows are understood to be driven by
powerful stellar winds (Masson & Chernin, 1993). Such winds are often detected
as blue-shifted absorption components in forbidden emission lines or in the
HeI 1.0830 $\mu$m line associated with protostars and young T Tauri stars
(e.g., Edwards et al., 1993; Kwan et al., 2007). Winds are also believed to
cause the shocked IR H2 emission seen in jets from heavily extinguished young
stellar objects. Such stellar winds are inferred to be driven by mass
accretion onto young stars. Jets detected in NIR H2 lines frequently are
displaced in space and in radial velocity from photospheric absorption
features and usually have relatively large linewidths of a couple dozen km s-1
or more. For example, in a recent VLT/ISAAC pilot study of the H2 1$-$0 S(1)
emission from embedded Class I sources, mean intensity-weighted velocities
were blue-shifted by -90 to -10 km s-1 and velocity widths of the lines varied
from $\sim$ 45 km s-1 to $\sim$ 80 km s-1 (Chrysostomou et al., 2008).

X-rays and UV may also excite molecular hydrogen emission in the inner
circumstellar environments of protostars. Both of these processes offer a
radiative alternative to the mechanical excitation excitation of the NIR H2
emission. If dominant, these excitation mechanisms may produce emission lines
with lower velocity widths or offsets than if the H2 were mechanically excited
in shocks or jets. We now consider the mechanisms, observational evidence, and
implications for these processes in protostellar environments.

It is now well established that Class I/FS protostars are copious X-ray
emitters (10${}^{29}\leq L_{x}\leq 10^{31}$ erg s-1, $\sim$100$-$1000 times
more X-ray luminous than main-sequence stars), and emit higher energy X-rays
(4-6 keV vs. 1-2 keV) than their older, T Tauri cousins (e.g., Casanova et
al., 1995; Grosso et al., 1997; Imanishi et al., 2001; Feigelson et al.,
2007). One would, therefore, expect significant X-ray heating of the
disk/inner envelope material surrounding these central objects (e.g.,
Meijerink et al., 2009, 2007). X-rays from young stellar objects (YSOs) can
penetrate disk atmospheres to fairly large surface densities and can ionize
circumstellar gas at a level greater than Galactic cosmic rays out to large
distances ($\sim$104 AU; Glassgold et al., 2005).

The mechanism of H2 excitation by X-rays requires impacts from energetic
electrons. X-rays impact hydrogen molecules, ejecting electrons. These high
energy electrons subsequently collide with and ionize or dissociate ambient
gas, losing kinetic energy in the process. Some of electrons will eventually
have energies appropriate to excite ambient H2 molecules into excited states
instead of dissociating them completely (e.g., Gredel & Dalgarno, 1995; Tine
et al., 1997; Maloney et al., 1996). Direct evidence for X-ray heating of gas
other than H2 in disks has been found for several Class I sources in which the
6.4 keV line from neutral iron has been detected (Giardino et al., 2007, and
references therein).

Glassgold and co-workers especially have emphasized the importance of X-ray
heating of disk atmospheres out to relatively large distances from the central
source (Glassgold et al., 2007, 2004). Such heating can extend to large
distances because of disk flaring, first proposed to account for the spectral
energy distributions (SEDs) produced by dust (Kenyon & Hartmann, 1987; Chiang
& Goldreich, 1997). More recent disk modeling, especially studies involving
disk chemistry, routinely use vertically stratified models, with different gas
and dust scale heights (d’Alessio et al., 1999; Aikawa et al., 2002; Gorti &
Hollenbach, 2008; Lacy et al., 2010). In X-ray heated disk models, there is a
low column density (N${}_{H}\sim$1020 cm-2) surface layer of hot (T$\sim$4000
K) gas that extends to $\gtrsim$10 AU radius (Najita et al., 2009).

In the context of disk heating by X-rays, it must be noted that current disk
models assume irradiation by a central source with X-ray spectra typical of T
Tauri stars: with plasma having $kT_{X}=$1 keV and a low-energy cut-off of 100
eV (e.g., Glassgold et al., 2007). However, Class I and FS sources are known
to have harder X-ray spectra, with 4 keV $\leq kT_{x}\leq$6 keV being typical
(e.g., Imanishi et al., 2001). Unsurprisingly, the column densities of
hydrogen gas inferred towards Class I/FS sources from X-ray observations are
$\sim\ N_{H}\ =\ 1-5\times 10^{22}$ cm-2, 1-2 orders of magnitude greater than
inferred for T Tauri stars. Finally, the quantity, Lx/Lbol, is systematically
smaller for Class I/FS objects relative to T Tauri stars, consistent with the
interpretation of higher accretion rates in these systems. The higher
accretion rates create higher optical depths in the X-ray absorbing gas,
obscuring the lower energy X-rays and producing relatively lower Lx/Lbol
ratios. Higher accretion rates would also lead to higher heating rates of the
disk gas in self-embedded protostars than in classical T Tauri star (CTTS)
disks (e.g., d’Alessio et al., 2004).

The possibility of UV excitation of the NIR H2 lines also exists for Class
I/FS protostars, given their consistently higher accretion rates relative to T
Tauri stars. The excess UV continuum emission observed in CTTSs has been
modeled as being produced by the impact of accretion columns onto the pre-
main-sequence stellar surface (e.g., Gullbring et al., 2000). This UV
continuum excess could potentially excite molecular hydrogen, producing
emission lines in the near-infrared. Strong Lyman-$\alpha$ emission from the
central object can irradiate the disk’s surface, and, if H2 is present at
$\sim$ 2000K, can excite the H2 into electronic states which produce a rich UV
emission-line spectrum as observed in some T Tauri stars (Herczeg et al.,
2006). Finally, if the stellar EUV flux is sufficiently strong, it can ionize
hydrogen, which produces high temperatures (T$\approx$ 104K) and small mean
molecular weights at the disk surface. Outside some critical radius, the gas
becomes unbound, and a slow, v$\approx$10 km s-1, photoevaporative wind forms
(Alexander et al., 2006a, b; Gorti et al., 2009; Woitke et al., 2009).

Like molecular hydrogen, [NeII] emission can also arise from high energy
photons in dense circumstellar disks or outflows. [NeII] line emission at
12.81 $\mu$m and [NeIII] emission at 15.5 $\mu$m was predicted to be
detectable in the case of X-ray heated disk gas. The 12.81 $\mu$m [NeII] line
has subsequently been detected in a number of young stars with the Spitzer
Space Telescope (Glassgold et al., 2007; Pascucci et al., 2007; Lahuis et al.,
2007; Flaccomio et al., 2009; Najita et al., 2010) and has been studied at
high spectral resolution from the ground in three objects (Herczeg et al.,
2007; Najita et al., 2009). Although a disk origin has been postulated or
confirmed for the [NeII] emission observed in the objects studied at high
spectral resolution, there are also sources in which [NeII] is detected in
outflows close to the sources (e.g., Neufeld et al., 2006; van Boekel et al.,
2009).

In general, gas motions in the inner circumstellar environments (within $\sim$
100 AU) of protostars have not been studied well except for a very few cases
of velocity resolved NIR CO observations, and this has limited our
understanding of how they accrete matter, form winds, shed angular momentum,
and disperse their natal circumstellar envelopes. We seek to understand these
processes as well as the UV and X-ray radiation environments of their
circumstellar disks by embarking on a new study of protostars’ NIR H2 line
strengths, morphologies, and velocities.

H2 NIR ro-vibrational line ratios yield excitation temperatures and also
provide clues to excitation mechanisms. The intensity ratio of $S(1)$ lines in
the $v\ =\ 2\rightarrow 1$ and $v\ =\ 1\rightarrow 0$ transitions is often
used as a diagnostic. This ratio has a value of 0.13 in shocked gas at 2000 K
and a value of 0.54 for UV pumping. For X-ray excitation, this ratio is
predicted to be 0.06 in gas of low fractional ionization (10-4) and 0.54 in
gas of high fractional ionization (10-2) (Gredel & Dalgarno, 1995). Any
observed variability in these emission lines can also provide clues to the
nature of their excitation and the location of the excited and emitting gas.

Several recent studies have made good progress in diagnosing the nature of H2
excitation in the inner circumstellar disks of CTTSs using observations of
their NIR ro-vibrational emission lines. A growing body of work shows that the
NIR H2 line radial velocities and linewidths of many CTTS and Herbig AeBe
stars are consistent with UV or X-ray excitation in a circumstellar disk
(Weintraub et al., 2000; Bary et al., 2002, 2003; Carmona et al., 2007; Bary
et al., 2008). By contrast, in a recently completed NIR adaptive optics (AO)
integral field spectroscopic study of several CTTS, the H2 emissions were most
consistent with shocks arising in winds (Beck et al., 2008).

Despite this progress in understanding H2 emission in CTTSs, little has been
done to date in probing the nature of such emission in more embedded Class I
and flat–spectrum protostars. Greene & Lada (1996) reported that these
protostars were significantly more likely to have NIR H2 emission than CTTSs,
and Doppmann et al. (2005) found that 23 of 52 observed protostars showed NIR
H2 emission. The most embedded protostars also have significant envelopes that
may have enough column density to generate observable H2 emission if the
hydrogen molecules there receive sufficient radiative or mechanical energy. If
this envelope source existed, it would be a new source of emission not present
in CTTSs. Unfortunately, earlier studies did not have adequate spectral
resolution or spectral range to observe multiple NIR ro–vibrational lines
simultaneously with high spectral resolution and good signal-to-noise: all
necessary ingredients for measuring line velocities and line ratios in order
to diagnose excitation.

An early motivation for a new observational study was the identification of
outflow drivers amongst late-stage protostars via detection of NIR H2 emission
at high spectral resolution observed directly towards the putative powering
source (Barsony, 2005). However, careful examination of the H2 $v=1-0$ $S(0)$
line profiles (at 2.2235 $\mu$m) observed at R$\sim$17,000 in the work of
Doppmann et al. (2005), showed that a substantial number of sources exhibited
relatively narrow linewidths ($\Delta v\leq$ 15 km s-1). Furthermore, in many
cases, the H2 emission line centers were not significantly displaced from the
central object’s radial velocity. Taken together, these two observations call
into question an outflow origin for the observed NIR H2 emission in some
sources.

We have conducted a new study of Class I and flat-spectrum protostars with
data sufficient for diagnosing the natures of their H2 emission in their inner
circumstellar disks and envelopes. We report on the sample and observations in
§2 and present our analysis of these data in §3. We discuss our results in §4
and summarize our work in §5.

##  2 Observations and Data Reduction

High resolution NIR spectra of 18 Class I and flat-spectrum protostars
previously searched for H2 $v=1-0$ $S(0)$ emission by Doppmann et al. (2005)
were re-observed with the Keck II telescope on Mauna Kea, Hawaii using its
NIRSPEC multi-order cryogenic echelle facility spectrograph (McLean et al.,
1998). This included 17 objects found to exhibit H2 $v=1-0$ $S(0)$ emission in
this previous study, and one (03260+311A) without. All new spectra were
acquired on 2007 June 24 and 25 UT ($\rho$ Oph and Ser objects) and 2008
January 24 and 25 (Tau-Aur and Per objects). The late type (K1 – M2.5) dwarfs
HD 20165, HD 28343, and HD 285968 with precision velocities measured by
Nidever et al. (2002) were observed on 2008 January 24 UT to serve as radial
velocity references.

Spectra were acquired with a 0$\farcs$58 (4-pixel) wide slit, providing
spectroscopic resolution $R\equiv\lambda/\delta\lambda$ = 18,000 (16.7 km
s-1). The plate scale was 0$\farcs$20 pixel-1 along the 12$\arcsec$ slit
length, and the seeing was typically 0$\farcs$5–0$\farcs$6\. The NIRSPEC
gratings were oriented to observe the 2.1218 $\mu$m H2 $v=1-0$ $S(1)$, 2.2477
$\mu$m H2 $v=2-1$ $S(1)$, and 2.386 $\mu$m H2 $v=3-2$ $S(1)$ lines on the
instrument’s 1024 $\times$ 1024 pixel InSb detector array in a single
exposure. The 2.2233 $\mu$m H2 $v=1-0$ $S(0)$ line previously observed by
Doppmann et al. (2005) was also captured in this grating setting. The
NIRSPEC-7 blocking filter was used to image these orders on the detector.
NIRSPEC was configured to acquire simultaneously multiple cross-dispersed
echelle orders 32–37 (2.05–2.40 $\mu$m, non-continuous) for all objects. Each
order had an observed spectral range $\Delta\lambda\simeq\lambda/67$ ($\Delta
v\simeq$ 4450 km s-1).

The slit was held physically stationary during the exposures and thus rotated
on the sky as the non-equatorially-mounted telescope tracked when observing.
Data were acquired in pairs of exposures of durations from 180–600 s each,
with the telescope nodded 3$\arcsec$ or 6$\arcsec$ along the slit between
frames so that object spectra were acquired in all exposures. Most of the
targets were observed twice on consecutive nights. The observation dates,
total integration times, slit angles and coordinates of all Class I and flat-
spectrum objects are given in Table 1. The early-type (B9–A0) dwarfs HD 28354,
HR 6070, HR 5993, and HD 168966 were observed for telluric correction of the
target spectra. The telescope was automatically guided with frequent images
from the NIRSPEC internal “SCAM” IR camera during all exposures of more than
several seconds duration. Spectra of the internal NIRSPEC continuum lamp were
taken for flat fields, and exposures of the Ar, Ne, Kr, and Xe lamps were used
for wavelength calibrations.

All data were reduced with IRAF. First, object and sky frames were differenced
and then divided by normalized flat fields. Next, bad pixels were fixed via
interpolation, and spectra were extracted with the APALL task. Spectra were
wavelength calibrated using low-order fits to lines in the arc lamp exposures,
and spectra at each slit position of each object were co-added. Instrumental
and atmospheric features were removed by dividing wavelength-calibrated object
spectra by spectra of early-type stars observed at similar airmass at each
slit position. Final spectra were produced for each night by combining the
spectra of both slit positions for each object and then multiplying them by
spectra of 10,000 K blackbodies to rectify the spectral shapes induced when
dividing by the telluric stars with that effective temperature. Spectra
acquired on different nights were not combined in any way.

##  3 Analysis and Results

The spectra of all objects in Table 1 were analyzed by measuring their H2
$v=1-0$ $S(1)$, $v=2-1$ $S(1)$, and $v=1-0$ $S(0)$ line fluxes, H2 line radial
velocities, and the radial velocities of photospheric absorption lines. The
2.386 $\mu$m H2 $v=3-2$ $S(1)$ line is in a region of poor atmospheric
transmission and was not significantly detected in any object. The values of
the H2 line properties and their night-to-night variations were analyzed for
physical insights into their excitation as described in this section.

Spectra of the H2 $v=1-0$ $S(1)$ region are shown for all objects in Figure 1;
the first epoch (2007 Jun 24 and 2008 Jan 24) is shown in the left panel and
the second (2007 Jun 25 and 2008 Jan 25) in the right. Of the 18 objects, 14
were observed in both epochs. When present, the $v=1-0$ $S(1)$ line was the
strongest H2 feature observed in all objects. All emission spectra shown in
the figures and all derived strength and velocity values presented in Table 2
are for the H2 emission that was spatially coincident with each object’s
continuum source, typically limited by the $0\farcs 6$ seeing, corresponding
to a spatial extent of $\sim 100$ AU for most sources.

###  3.1 H2 Emission Morphologies

The spatial extent of the H2 emission along the spectrograph slit was measured
for each object by examining the differenced and flat-fielded spectral images
that were created prior to spectral extraction. The non-rotating spectrograph
slit was projected onto the sky at different position angles each night, so
the two observational epochs sample the spatial extent of any extended
emission differently (see position angles in Table 1).

The angular extent of the H2 emission along the length of the slit is reported
in Table 2) for each observation. Ten of the 18 protostars showed H2 emission
that was over $1\arcsec$ in spatial extent in at least one observation (at
least one position angle). Of the 14 objects observed in two epochs, 6 had H2
$v=1-0$ $S(1)$ emission extended by $\sim 1\arcsec$ or less on both nights.
Therefore the molecular hydrogen emission of these 6 objects is contained
within about 70 AU ($\sim 0\farcs 5$) of their central stars. The highest H2
surface brightness was generally found to be coincident with each object’s
unresolved point source, but in some cases the extended emission may have more
integrated flux and luminosity than the point source component. Beck et al.
(2008) found this to be true in their AO integral field study of several CTTSs
that sampled their immediate circumstellar environments completely instead of
evaluating just two slit position angles as done in this study.

Objects with significantly different spatial extents between epochs (position
angles) may have their H2 emission confined to asymmetrical structures like
jets. 03260+3111B, 04158+2805, and 04264+2433 are examples of objects with
these potentially jet-like morphologies. 03260+3111A shows significantly
extended emission, but this is spatially displaced from the stellar continuum
which has no coincident H2 emission.

###  3.2 Variability

We found H2 emission in all of the 17 objects also found to have H2 emission
at earlier epochs by Doppmann et al. (2005). The one source without emission
(03260+3111A) in that earlier study also did not show H2 emission in our new
observations. This indicates that these these protostars do not terminate or
initiate their H2 emissions on time scales as short as 5 – 10 years,
suggesting that their H2 lines are emitted over physically large regions or
else are excited by processes with relatively stable fluxes.

Fourteen of the 18 objects were observed twice, with the two observations
spaced approximately 24 hr apart. The position angles of both observations
were sometimes similar (within $\sim 10\deg$), but often they varied by
$60\deg$ or more (Table 1). Therefore the different slit position angles must
be considered when comparing H2 line fluxes measured on the different dates to
determine whether any differences are due to true temporal variation or the
rotation of different spatial features (i.e., jets) onto or off of the slit.
However, the measured equivalent H2 emission widths and line velocities were
similar for both observations of each object (see Table 2), so there appears
to be little temporal variation on $\sim 1$ day time scales for the emission
that is spatially confined to be coincident with the protostar continuum
source confined within the slit width.

###  3.3 Radial Velocities and Line Widths

Stellar photospheric radial velocities were computed from the continuum
absorption lines of all observed objects. First, the object spectra were
cross-correlated (using fxcor in IRAF) with spectra of the radial velocity
standards, always HD 28343 (K7 V) and sometimes HD 20165 (K1 V) or HD 285968
(M2.5 V), using up to 4 spectral orders that contained photospheric lines but
no emission lines. Heliocentric $V_{LSR}$ radial velocities were computed for
each object by summing the radial velocity shift measured with the cross-
correlations, the mean measured radial velocities of the 3 standards (Nidever
et al., 2002), and the $V_{LSR}$ corrections for the objects and standards at
their time of observation (from the IRAF rvcorrect task). The $V_{LSR}$ radial
velocities computed in this way generally agreed well with those of the 10
objects also reported by Covey et al. (2006). We observed these 10 objects a
total of 18 times, and we measured the offset from the Covey et al. (2006)
velocities to be $2.2\pm 3.5$ km s-1. This offset is consistent with zero, and
the standard deviation is not surprising given the 17 km s-1 spectral
resolution and the generally heavily veiled spectra.

Radial velocities of H2 lines were computed by measuring the central
wavelengths of Gaussian profiles fit to the lines using the IRAF splot task
and then converting these values to velocities. These observed radial
velocities were converted to $V_{LSR}$ values by adding the $V_{LSR}$
correction offset computed for each object as described above. The velocities
of the H2 lines relative to photospheric lines of each object were computed by
differencing these two $V_{LSR}$ radial velocities, and the results are shown
in Table 2 and Figure 2. FWHM velocity widths of the $v=1-0$ $S(1)$ lines were
computed by subtracting the instrumental line width of 17 km s-1 in quadrature
from the FWHM values of the Gaussian fits. These resultant line widths are
also reported in Table 2, and their histogram is shown in Figure 3. Objects
observed twice on successive nights had similar radial and FWHM velocities
(Table 2), so these values were averaged to reduce noise in Figures 2 and 3.
The other H2 lines generally had similar FWHM values, but the $v=1-0$ $S(1)$
lines had the highest signal-to-noise, so only those values are presented. We
estimate that all reported velocities have uncertainties of a few km s-1.

###  3.4 H2 Line Fluxes and Ratios

Line luminosities were estimated by scaling the relative fluxes measured in
each H2 line to the spatially coincident 2.2 $\mu$m continuum and multiplying
this by the absolute 2.2 $\mu$m continuum flux estimated from each object’s
2MASS K-band magnitude after correcting for extinction. Extinctions were
calculated by de-reddening each object’s $JHK$ 2MASS magnitudes to the CTTS
locus (Meyer et al., 1997). Extinctions were estimated at each H2 line
wavelength using $A_{v}=9.09[(J-H)-(J-H)_{0}]$, $A_{k}=0.09A_{v}$, and
$A_{\lambda}\propto\lambda^{-1.9}$. These values were computed by and derived
from Cohen et al. (1981) for the CIT photometric system, which is essentially
identical to that of 2MASS (Carpenter, 2001).

Distances were assumed to be 140 pc for Tau-Aur (Kenyon et al., 1994), 140 pc
for $\rho$ Oph (Mamajek, 2008), 260 pc for SVS 2 in Serpens (Straižys et al.,
1996), and 320 pc for the 03260+3111 objects in Perseus (Herbig, 1998). H2
line equivalent widths and luminosities are presented in Table 2, and a
histogram of luminosities of the H2 $v=1-0$ $S(1)$ line is shown in Figure 4.
A histogram of H2 $1-0/2-1$ $S(1)$ line ratios is presented in Figure 5. As
done in previous figures, values derived from observations of objects acquired
on successive nights are averaged in these figures as well. Note that the H2
line ratio value plotted in Figure 5 is the inverse of the values presented in
Table 2 (column 8).

###  3.5 Correlations

We examined whether correlations exist between measured H2 line properties and
other protostellar activity indicators in an attempt to isolate the origins of
the H2 line emissions.

First we computed the correlations between the H2 line properties measured in
this new survey. The FWHM velocity widths and the $1-0/2-1$ $S(1)$ line ratios
in Table 2 have a correlation coefficient of 0.61, which improves to 0.71 when
both observations of a single object are averaged into single points. This
value drops to only 0.17 (individual or mean values) if the observations of
03260+3111B are excluded; its high FWHM velocity and relatively high $1-0/2-1$
$S(1)$ line ratio drives this correlation. Thus the object sample as a whole
does not show a good correlation between its H2 FWHM velocities and $1-0/2-1$
$S(1)$ line ratios. We then examined correlations between H2 FWHM velocities
and velocity offsets between H2 and photospheric lines. These values had
correlation coefficients of -0.43 and -0.47 for individual and averaged
values, respectively. The square of the correlation coefficient is below 0.25
in both cases, indicating that less than 25% of the variance of the two
quantities are in common for this sample of protostars, a poor correlation.

Next we correlated the protostars’ H2 line properties with other physical
characteristics measured in other studies. Twelve of the 18 objects have
measured X-ray luminosities or upper limits (Güdel et al., 2007; Flaccomio et
al., 2009), and the correlation coefficient between the logarithms of their H2
$1-0$ $S(1)$ line luminosities and the logarithms of their X-ray luminosities
is -0.20, suggesting a very weak or nonexistent inverse correlation between
these properties. Finally, we re-analyzed the spectra of Doppmann et al.
(2005) and used their equivalent width measurements of H2 $1-0$ $S(0)$ and HI
Br $\gamma$ emission lines to evaluate the correlation of these properties in
that somewhat earlier epoch. We found that these values had a correlation
coefficient of 0.20, indicating another poor correlation. In summary, we find
little correlation among the NIR H2 line emission properties or between these
properties and other young stellar activity indicators.

##  4 Discussion

The different radiative and collisional excitation mechanisms of H2 are well
matched to the radiative and mass flux processes in the environments of
protostars and T Tauri stars. We now interpret the results of the preceding
analysis in terms of several of these possible processes in order to constrain
the H2 excitation mechanisms of the sample and to understand better the
circumstellar environments of these protostars.

###  4.1 Emission Morphologies, Variability, and Velocities

In addition to their compact H2 line emissions, all but 4 of the 18 objects
also showed $v=1-0$ $S(1)$ molecular hydrogen emission extended by $\sim
1\arcsec$ or more along the slit (see Table 2), corresponding to 70 AU ($\sim
0\farcs 5$) or more projected radial distance. It is unlikely that UV
radiation could travel that far from the central protostars without
significant attenuation by gaseous and dusty envelopes, so it is likely that
this extended emission is excited by either stellar winds or high energy
X-rays. The expected small size of the UV emission region on the protostellar
photosphere also suggests that UV may not be a good candidate for exciting the
observed steady and long-lived H2 line emissions. Two objects, GY 21 and IRS
43, have extended $1-0$ $S(1)$ emission with broad line widths, FWHM $\gtrsim$
40 km s-1, about twice that of the compact emission spatially coincident with
their stellar continua. The extended emissions of these 2 objects are good
candidates for excitation in shocks caused by stellar winds.

The relatively stable values of the point source H2 emission over 1 day and
several-year time scales (see §3.2 and Table 1) also provide clues to the
nature of these emissions. Numerous Class I and FS protostars have been
observed to undergo rapid (several hr), large amplitude X-ray emission
variations (e.g., Imanishi et al., 2001, 2003; Güdel et al., 2007; Flaccomio
et al., 2009). This X-ray flaring of several protostars (e.g., IRS 43) has
also been observed to appear or disappear in data taken $\sim$ 5 – 10 years
apart. If these X-ray flares were exciting H2 close to the stars, then it is
likely that we would see significant night-to-night or year-to year variations
in their point-source near-IR line fluxes, but this is not seen in our data.
The very stable observed molecular hydrogen emission is more consistent with
mechanical (wind or jet) excitation as well as excitation by steady, non-
flaring X-ray emission from the protostars.

The preceding analysis of the velocity widths and velocity shifts of the H2
line emission in §3.3 also provides clues to the nature of its excitation.
Collisional excitations in jets or winds are likely to result in H2 emission
line radial velocities displaced from photospheric absorption lines by over 10
km s-1, exhibition of broad line wings, or large full-width half maximum
velocities of several 10 km s-1 or more (Maloney et al., 1996; Montmerle et
al., 2000; Nomura et al., 2007; Beck et al., 2008). Jets are also often
significantly collimated and spatially extended, making them easy to identify
in one or two dimensional spectral images (e.g., Schwartz & Greene, 2003).

The significant spatial extension of the H2 line emission of many objects is
consistent with collisional excitation in winds or by jets. However, only 5 of
13 protostars show radial velocity offsets between their H2 lines and
photospheric absorption lines with absolute value of greater than 4 km s-1
(see Fig. 2). This value is similar to the 3.5 km s-1 uncertainty we measured
for radial velocity standards (see §3.3), so we do not consider velocity
offsets less than 4 km s-1 to be significant. Only 3 protostars have velocity
offsets of at least 10 km s-1, significant at about the 3-$\sigma$ confidence
level or greater. This evidence suggests that collisional excitation in jets
is unlikely to be the molecular hydrogen excitation mechanism in most objects
(except the 3 with significant velocity differences).

However, the on-source H2 $v=1-0$ $S(1)$ line FWHM line widths of all 17
protostars with this on-source feature are broader than 10 km s-1 (see Fig.
3). Six of the emitting objects (35%) exhibit FWHM greater that 20 km s-1. All
6 of these objects also show spatially extended H2 emission many tens of AU
from their central stars; this combination of factors makes them good
candidates for collisional excitation in jets or winds. Seven of the emitting
objects (41%) have FWHM line widths below 16 km s-1. This is similar to the 9
to 14 km s-1 line widths found by Bary et al. (2003) and Bary et al. (2008)
for 8 of 10 T Tauri stars found to have H2 emission, which they interpreted as
evidence for quiescent emission in circumstellar disks. We conclude that the
molecular hydrogen emission line velocities and FWHM values of at least 3 to 6
of the 17 emitting objects are consistent with collisional excitation in jets,
and at least 7 or 8 objects have H2 velocity parameters consistent with
quiescent (non-collisional) excitation.

###  4.2 Emission Line Strengths and Ratios

NIR vibrational H2 line ratios are also sensitive to the gas excitation levels
and excitation mechanisms. Gredel & Dalgarno (1995) show that the ratios of
the H2 $v=1-0$ $S(1)$ to $v=2-1$ $S(1)$ lines are relatively sensitive to
excitation mechanisms. They compute the ratios of these 2 lines to be 1.9 for
UV excitation, 7.7 for shocked gas at $T=2000$ K, and 16.7 for X-ray
excitation of low ionization H2. Black & van Dishoeck (1987) also found
similar differences between UV and shock excitation of H2. However, there are
limits to the usefulness of these ratios as diagnostics of excitation in
circumstellar disks. In practice it is difficult to distinguish between
shocked and X-ray excited H2 emission from examining only a few NIR lines.
Collisions will thermalize the excitation levels of H2 in a sufficiently dense
gas, so $v=1-0$ $S(1)$ to $v=2-1S(1)$ line ratios indicative of cold-to-warm
gas in equilibrium are not always useful for distinguishing between excitation
mechanisms (Gredel & Dalgarno, 1995; Maloney et al., 1996; Tine et al., 1997;
Nomura et al., 2007; Beck et al., 2008).

However, these line ratios may be useful for diagnosing excitation processes
in extreme cases. Black & van Dishoeck (1987) note that fluorescent UV
excitation of H2 produces significant population of vibrational levels $v\geq
2$ and therefore strong emission in the $v=2-1$ $S(1)$ line when not
thermalized in a high density environment. This level population can be
characterized by a temperature of $T_{\rm vib}\approx 6000-9000$ K, much
larger than the $T_{\rm vib}=T\simeq 2000$ K characteristic of shock
excitation. Therefore any objects with observed ratios of H2 $v=1-0$ $S(1)$ to
$v=2-1$ $S(1)\simeq 2$ ($2-1/1-0\simeq 0.5$) may be exhibiting UV-excited H2
emission. However, there are no objects in our sample with H2 $v=2-1/1-0$
$S(1)>0.25$, so we do not have any good candidates for purely UV excitation in
a dust-free environment as modeled by Black & van Dishoeck (1987).

The presence of dust grains can significantly alter these ratios and
complicate their interpretation. Nomura et al. (2007) have computed the
expected NIR H2 line fluxes for conditions in circumstellar disks around young
stars, modeling X-ray and UV heating in the presence of both gas and dust and
accounting for thermalization at high gas densities. They find that the H2
$v=2-1$ $S(1)$ to $v=1-0$ $S(1)$ line ratio is greatly impacted by the
presence of dust grains of different sizes (see their Figure 17). For UV
excited emission, Nomura et al. (2007) find that this line ratio is $\simeq
0.025$ for a power law dust grain distribution with a maximum size of 10
$\mu$m - 1 mm, roughly consistent with that expected for a protostar’s
circumstellar disk. They find that the grains must be much larger ($\sim 10$
cm or more) for this ratio to approach the dust-free value of 0.5 computed by
Gredel & Dalgarno (1995). This line ratio is much less sensitive to grain size
in the case of X-ray excitation; Nomura et al. (2007) find that this value is
close to the Gredel & Dalgarno (1995) value of 0.06 for a maximum grain size
of 10 $\mu$m - 10 cm. Thus they find that the H2 $v=2-1$ $S(1)$ to $v=1-0$
$S(1)$ line ratio differs by only about a factor of 2 for a circumstellar disk
with a power law grain size distribution with a maximum size 10 $\mu$m - 1 mm.
Nomura et al. (2007) do not consider collisional excitation by winds or jets,
but collisional excitation may produce a fairly wide range of gas temperatures
and H2 line ratios as discussed previously.

If excited by collisions in shocks, ratios of the $v=1-0$ $S(1)$ to $v=1-0$
$S(0)$ emission lines can be used to estimate ortho:para ratios of molecular
hydrogen and to assess the thermal history of the emitting gas. The values of
ortho:para ratios were modeled for C- and J-type shocks by Wilgenbus et al.
(2000); see also their summary of previous work. Kristensen et al. (2007) and
Harrison et al. (1998) showed that the $v=1-0$ $S(1)$ to $v=1-0$ $S(0)$
emission line ratio directly yields the molecular hydrogen ortho:para ratio
with little sensitivity to H2 rotation temperature. Using Eq. 5 of Kristensen
et al. (2007) and assuming an H2 rotation temperature of 3500 K, we find that
the mean ortho:para ratio for our object sample is $<o/p>=3.2\pm 0.8$.
Hydrogen atom exchanges in shocks set the high temperature limit to o/p $\leq
3$ (e.g., see Wilgenbus et al., 2000). Thus it appears that not all objects in
our sample have molecular hydrogen emission consistent with production in
shocks since a number of objects have o/p $>3$. Unfortunately we were unable
to use line ratios to diagnose the nature of the spatially extended H2
emissions seen in many objects (see Table 2 and §4.1). This emission was
generally much weaker than the point source emission, and it was not detected
significantly in any NIR H2 line except $v=1-0$ $S(1)$ for any object.

###  4.3 Molecular Hydrogen Excitation Mechanisms in Observed Protostars

Our sample has 5 objects with H2 $v=2-0$ $S(1)$ to $v=1-0$ $S(1)$ line ratios
in the 0.025 - 0.06 range, consistent with UV or X-ray excitation in the
presence of dust. Of these, 04264+2433, 04295+2251, 04365+2535, and IRS 43
have relatively low mean H2 $v=1-0$ $S(1)$ line widths, FWHM $\lesssim 15$ km
s-1. All but 04264+2433 have been detected in X-rays, so X-ray or UV
excitation may be possible for these protostars. WL 12 is the other protostar
with a line ratio in this range, and it has a broad FWHM $\simeq 30$ km s-1
and is associated with a molecular outflow (Bontemps et al., 1996). Therefore
its H2 may be collisionally excited. However, WL 12 has ortho:para ratios of
about 4.5, larger than the o/p $\leq 3$ limit that can be produced in C-shocks
or J-shocks. The objects 04158+2805, 04181+2654, and WL 6 also have estimated
o/p $\geq 3.5$, indicating non-shock excitation even if their H2 rotation
temperature is somewhat higher than the assumed 3500 K. Therefore the NIR H2
emission of these objects may not be excited in jets. The latter 3 objects
also have H2 radial velocity offset by less than 5 km s-1 from their
photospheric velocities. However, 04158+2805 and 04181+2654 have H2 FWHM line
widths $\geq$ 20 km s-1, clouding a non-mechanical interpretation of their
molecular hydrogen excitation.

Interestingly, there are several objects in the sample whose observed
velocities and line fluxes suggest quiescent, non-mechanical origins for their
molecular hydrogen emissions. 04361+2547, WL 6, and IRS 67 all have small H2
FWHM line widths, small H2 velocity offsets from photospheric velocities, and
small H2 emission spatial extents (See Table 2). Interestingly, at least one
measurement of the $v=1-0$ $S(1)$ to $v=2-1$ $S(1)$ line ratios is $\sim$0.07
for each object, similar to the value of 0.06 computed by Gredel & Dalgarno
(1995) and Nomura et al. (2007) for X-ray excitation of H2. Thus these objects
appear to be the best candidates for non-mechanical excitation of their
molecular hydrogen emissions.

We conclude this discussion by noting that several of the protostars have NIR
ro-vibrational emission properties consistent with collisional excitation, and
some others appear to be good candidates for X-ray and / or UV excitation.
This is similar to the results found by Beck et al. (2008) and Bary et al.
(2008) in their surveys of T Tauri stars. We also find no individual or set of
spectral features that are inconsistent with previous observations of NIR H2
emission in CTTSs. It appears that there is no single NIR line diagnostic that
can clearly identify the excitation mechanisms of H2 in the circumstellar
disks and environments of protostars, and correlations between diagnostics are
not strong (§3.5). However, the measures of emission line morphologies,
velocity widths, velocity shifts, and line ratios can constrain the various
emission mechanisms when interpreted within an appropriate theoretical model.

##  5 Summary

We present new observations of near-infrared H2 line emission in a sample of
18 Class I and flat-spectrum low mass protostars, primarily in the Tau-Aur and
$\rho$ Oph dark clouds. We reach the following conclusions from analyzing
these data:

1\. All 17 objects found to have NIR H2 ro-vibrational line emission spatially
conincident with their continuum sources in an earlier epoch were also found
to have this emission in this new study, 5 – 10 years later. There appears to
be little temporal variation of this emission on $\sim 1$ day time scales. Ten
of the 18 protostars showed H2 $v=1-0$ $S(1)$ line emission that was over
$1\arcsec$ in spatial extent in at least one observation (at least one
position angle).

2\. Nearly all of the protostars have H2 $v=1-0$ $S(1)$ line emission radial
velocities within 10 km s-1 of their stellar photospheric line velocities;
only 3 objects have H2 velocity offsets greater than or equal to 10 km s-1.
This evidence suggests that collisional excitation in jets is unlikely to be
the molecular hydrogen excitation mechanism in many objects.

3\. The H2 $v=1-0$ $S(1)$ line FWHM line widths of all 17 protostars with this
feature on-source are broader than 10 km s-1 (Fig. 3). Six of the emitting
objects (35%) exhibit FWHM greater that 20 km s-1, and these are good
candidates for collisional excitation in jets or winds. Seven of the emitting
objects (41%) have point source FWHM line widths below 16 km s-1. The
spatially extended H2 $v=1-0$ $S(1)$ line emission of two objects had line
widths FWHM $\gtrsim$ 40 km s-1, about twice that of their central point
source emission. This is consistent with collisional excitation by jets or
winds.

4\. The molecular hydrogen emission line velocities and FWHM values of at
least 3 to 6 of the 17 objects with on-source emission are consistent with
collisional excitation in jets. At least 7 or 8 objects have H2 velocity
parameters consistent with quiescent (non-collisional) excitation. There are
several objects whose small emission line widths, small H2 – photospheric
radial velocity differences, and small spatial extents are more consistent
with quiescent molecular hydrogen emission and not collisional excitation.

5\. Several of the protostars have H2 $v=2-0$ $S(1)$ to $v=1-0$ $S(1)$ line
ratios indicative of X-ray or UV excitation (in the presence of dust) and are
known X-ray emitters. 04361+2547, WL 6, and IRS 67 are the best examples of
such protostars. However, we see no rapid variation in the H2 $\Delta v=1$
$S(1)$ line fluxes on $\sim$24 hr time scales as might be expected from
excitation by X-ray flaring events.

6\. We find that the mean ortho:para ratio for our object sample is
$<o/p>=3.2\pm 0.8$. Hydrogen atom exchanges in shocks set the high temperature
limit to o/p $\leq 3$ (e.g., see Wilgenbus et al., 2000). Thus it appears that
not all objects in our sample have molecular hydrogen emission consistent with
production in shocks since a number of objects have o/p $>3$. 04158+2805,
04181+2654, WL 6, and WL 12 are all estimated to have ortho:para ratios
significantly higher than this value. However, WL 6 is the only protostar with
H2 line ratios and velocities also indicative of non-mechanical excitation.

We thank D. Hollenbach and U. Gorti for helpful discussions of our data and
its interpretation via theoretical models. We also thank G. Herczeg for
discussing pre-publication data and thank the anonymous referee for thoughtful
suggestions that improved this paper. The Keck Observatory Observing
Assistants H. Hershley and C. Parker are thankfully acknowledged for
assistance with the observations. The authors wish to recognize and
acknowledge the very significant cultural role and reverence that the summit
of Mauna Kea has always had within the indigenous Hawaiian community. We are
most fortunate to have the opportunity to conduct observations from this
mountain. TPG acknowledges support from NASA’s Origins of Solar Systems
program via WBS 811073.02.07.01.89. MB and TPG would like to acknowledge NASA
support via NExScI for travel expenses to the W.M. Keck Observatory for
acquiring the observations for this project. Facilities: Keck (NIRSpec), IRAF,
2MASS

## References

  * Aikawa et al. (2002) Aikawa, Y., van Zadelhoff, G.J., van Dishoeck, E.F., & Herbst, E. 2002, A&A, 386, 622
  * Alexander et al. (2006a) Alexander, R.D., Clarke, C.J., & Pringle, J.E. 2006a, MNRAS, 369, 216
  * Alexander et al. (2006b) Alexander, R.D., Clarke, C.J., & Pringle, J.E. 2006b, MNRAS, 369, 229
  * Bary et al. (2002) Bary, J. S., Weintraub, D. A., & Kastner, J. H. 2002, ApJ, 576, L73
  * Bary et al. (2003) Bary, J. S., Weintraub, D. A., & Kastner, J. H. 2003, ApJ, 586, 1136
  * Bary et al. (2008) Bary, J. S., Weintraub, D. A., Shukla, S., Leisenring, J.M., & Kastner, J. H. 2008, ApJ, 678, 1088
  * Barsony (2005) Barsony, M. 2005, in “High-Resolution Infrared Spectroscopy in Astronomy,” Proc. of an ESO Workshop held at Garching, Germany, 18-21 November 2003, eds. Käufl, H.U., Siebenmorgen, R., & A. Moorwood, pp.125-130
  * Beck et al. (2008) Beck, T. L., McGregor, P. J., Takami, M., & a, T.-S. 2008, ApJ, 676, 472
  * Black & van Dishoeck (1987) Black, J. H., & van Dishoeck, E. F. 1987, ApJ, 322, 412
  * Bontemps et al. (1996) Bontemps, S., Andre, P., Terebey, S., & Cabrit, S. 1996, A&A, 311, 858
  * Burton et al. (1989) Burton, M. G., Brand, P. W. J. L., a, T. R., & Webster, A. S. 1989, MNRAS, 236, 409
  * Carmona et al. (2007) Carmona, A., van den Ancker, M.E., Henning, Th., Goto, M., Fedele, D., & Stecklum, B. 2007, A&A, 476, 853
  * Carpenter (2001) Carpenter, J. M. 2001, AJ, 121, 2851
  * Casanova et al. (1995) Casanova, S., Montmerle, T., a, E. D., & Andre, P. 1995, ApJ, 439, 752
  * Chiang & Goldreich (1997) Chiang, E.I. & Goldreich, P. 1997, ApJ, 490, 368
  * Chrysostomou et al. (2008) Chrysostomou, A., Bacciotti, F., Nisini, B., Ray, T.P., Eislöffel, J., Davis, C.J., & Takami, M. 2008, A&A, 482, 575
  * Cohen et al. (1981) Cohen, J. G., a, S. E., Elias, J. H., & aa, J. A. 1981, ApJ, 249, 481
  * Covey et al. (2006) Covey, K. R., Greene, T. P., Doppmann, G. W., & a, C. J. 2006, AJ, 131, 512
  * d’Alessio et al. (1999) d’Alessio, P., Calvet, N., Hartmann, L., Lizano, S., & Cantó, J. 1999, ApJ, 527, 893
  * d’Alessio et al. (2004) d’Alessio, P., Calvet, N., Hartmann, L., Muzerolle, J., & Sitko, M. 2004, in Star Formation at High Angular Resolution, IAU Symp. 221, 403, eds. M.G. Burton, R. Jayawardhana, & T.L. Bourke
  * Davis et al. (2010) Davis, C.J., Gell, R., Khanzadyan, T., Smith, M.D., & Jenness, T. 2010, A&A, 511, 24
  * Doppmann et al. (2005) Doppmann, G. W., Greene, T. P., Covey, K. R., & Lada, C. J. 2005, AJ, 130, 1145
  * Edwards et al. (1993) Edwards, S., Ray, T., & Mundt, R. 1993, Protostars and Planets III, 567
  * Feigelson et al. (2007) Feigelson, E., Townsley, L., Güdel, M., & Stassun, K. 2007, Protostars and Planets V, 313
  * Flaccomio et al. (2009) Flaccomio, E., Stelzer, B., Sciortino, S., Micela, G., Pillitteri, I., & Testi, L. 2009, A&A, 505, 695
  * Giardino et al. (2007) Giardino, G., Favata, F., Pillitteri, I., Flaccomio, F., Micela, G., & Sciortino, S. 2007, A&A, 475, 891
  * Glassgold et al. (2005) Glassgold, A.E., Feigelson, E.D., Montmerle, T., & Wolk, S. 2005, ASP Conf. Ser. 341, p.165
  * Glassgold et al. (2007) Glassgold, A.E., Najita, J.R., Igea, J. 2007, ApJ, 656, 515
  * Glassgold et al. (2004) Glassgold, A.E., Najita, J.R., Igea, J. 2004, ApJ, 615, 972
  * Gorti & Hollenbach (2008) Gorti, U. & Hollenbach, D. 2008, ApJ, 683, 287
  * Gorti et al. (2009) Gorti, U., Dullemond, C.P., & Hollenbach, D. 2009, ApJ, 705, 1237
  * Gredel & Dalgarno (1995) Gredel, R., & a, A. 1995, ApJ, 446, 852
  * Greene & Lada (1996) Greene, T. P., & Lada, C. J. 1996, AJ, 112, 2184
  * Greene & Lada (2002) Greene, T. P., & Lada, C. J. 2002, AJ, 124, 2185
  * Grosso et al. (1997) Grosso, N., Montmerle, T., Feigelson, E. D., André, P., Casanova, S., & Gregorio-Hetem, J. 1997, Nature, 387, 56
  * Güdel et al. (2007) Güdel, M., et al. 2007, A&A, 468, 353
  * Gullbring et al. (2000) Gullbring, E., Calvet, N., Muzerolle, J., & Hartmann, L. 2000, ApJ, 544, 927
  * Harrison et al. (1998) Harrison, A., Puxley, P., Russell, A., & Brand, P. 1998, MNRAS, 297, 624
  * Herbig (1998) Herbig, G. H. 1998, ApJ, 497, 736
  * Herczeg et al. (2007) Herczeg, G. J., Najita, J. R., Hillenbrand, L. A., & Pascucci, I. 2007, ApJ, 670, 509
  * Herczeg et al. (2006) Herczeg, G. J., Linsky, J. L., Walter, F. M., Gahm, G. F., & Johns-Krull, C. M. 2006, ApJS, 165, 256
  * Imanishi et al. (2001) Imanishi, K., Koyama, K., & Tsuboi, Y. 2001, ApJ, 557, 747
  * Imanishi et al. (2003) Imanishi, K., Nakajima, H., Tsujimoto, M., Koyama, K., & Tsuboi, Y. 2003, PASJ, 55, 653
  * Kenyon et al. (1994) Kenyon, S. J., Dobrzycka, D., & Hartmann, L. 1994, AJ, 108, 1872
  * Kenyon & Hartmann (1987) Kenyon, S. J., & Hartmann, L. 1987, ApJ, 323 714
  * Kristensen et al. (2007) Kristensen, L. E., Ravkilde, T. L., Field, D., Lemaire, J. L., & Pineau Des Forêts, G. 2007, A&A, 469, 561
  * Kwan et al. (2007) Kwan, J., Edwards, S., & Fischer, W. 2007, ApJ, 657, 897
  * Lacy et al. (2010) Lacy, J.H., Harrold, S.T., Watson, D.M. 2010, BAAS, 41, 346
  * Lahuis et al. (2007) Lahuis, F., van Dishoeck, E.F., Blake, G.A., Evans, N.J. II, Kessler-Silacci, J.E., & Pontopiddan, K.M. 2007, ApJ, 665, 492L
  * Maloney et al. (1996) Maloney, P. R., Hollenbach, D. J., & Tielens, A. G. G. M. 1996, ApJ, 466, 561
  * Mamajek (2008) Mamajek, E. E. 2008, Astronomische Nachrichten, 329, 10
  * Masson & Chernin (1993) Masson, C. R., & Chernin, L. M. 1993, ApJ, 414, 230
  * McLean et al. (1998) McLean, I. S. et al. 1998, Proc. SPIE, 3354, 566
  * Meyer et al. (1997) Meyer, M. R., Calvet, N., & Hillenbrand, L. A. 1997, AJ, 114, 288
  * Montmerle et al. (2000) Montmerle, T., Grosso, N., Tsuboi, Y., & Koyama, K. 2000, ApJ, 532, 261
  * Moriarty-Schieven et al. (1992) Moriarty-Schieven, G. H., Wannier, P. G., Tamura, M., & Keene, J. 1992, ApJ, 400, 260
  * Meijerink et al. (2009) Meijerink, R., Poelman, D.R., Glassgold, A.E., Najita, J.R., Tielens, A.G.G.M., & Spaans, M. 2009, “The Evolving ISM in the Milky Way and Nearby Galaxies,” Proc. of the 4th Spitzer Science Center Conference, eds. K. Sheth, A. Noriega-Crespo, J. Ingalls, & R. Palladini, on-line at http://ssc.spitzer.caltech.edu/mtgs/ismevol/
  * Meijerink et al. (2007) Meijerink, R., Glassgold, A., & Najita, J. 2007, in “The Spirit of Bernard Lyot: The Direct Detection of Planets and Circumstellar Disks in the 21st Century,” Proc. of Conf. held 4-8 June 2007 at U.C. Berkeley, ed. P. Kalas
  * Najita et al. (2010) Najita, J.R., Carr, J.S., Strom, S.E., Watson, D.M., Pascucci, I., Hollenbach, D., Gorti, U., & L. Keller, 2010, ApJ, 712, 274
  * Najita et al. (2009) Najita, J.R., Doppmann, G.W., Bitner, M.A., Richter, M.J., Lacy, J.H., Jaffe, D.T., Carr, J.S., Meijerink, R., Blake, G.A., Herczeg, G.J.., & Glassgold, A.E. 2009, ApJ, 697, 957
  * Neufeld et al. (2006) Neufeld, D., et al. 2006, ApJ, 649, 816
  * Nidever et al. (2002) Nidever, D. L., Marcy, G. W., Butler, R. P., Fischer, D. A., & Vogt, S. S. 2002, ApJS, 141, 503
  * Nomura et al. (2007) Nomura, H., Aikawa, Y., Tsujimoto, M., Nakagawa, Y., & Millar, T. J. 2007, ApJ, 661, 334
  * Pascucci et al. (2007) Pascucci, I., Hollenbach, D., Najita, J., Muzerolle, J., Gorti, U., Herczeg, G.J., Hillenbrand, L.A., Kim, J.S., Carpenter, J.M., Meyer, M.R., Mamajek, E.E., & Bouwman, J. 2007, ApJ, 663, 383
  * Rieke & Lebofsky (1985) Rieke, G. H., & Lebofsky, M. J. 1985, ApJ, 288, 618
  * Schwartz & Greene (2003) Schwartz, R. D., & Greene, T. P. 2003, AJ, 126, 339
  * Straižys et al. (1996) Straižys, V., Černis, K., & Bartašiūtė, S. 1996, Baltic Astronomy, 5, 125
  * Tine et al. (1997) Tine, S., Lepp, S., Gredel, R., & Dalgarno, A. 1997, ApJ, 481, 282
  * van Boekel et al. (2009) van Boekel, R., Güdel, M., Henning, Th., Lahuis, F., & Pantin, E. 2009, A&A, 497, 137
  * Weintraub et al. (2000) Weintraub, D. A., Kastner, J. H., & Bary, J. S. 2000, ApJ, 541, 767
  * White & Hillenbrand (2004) White, R. J., & Hillenbrand, L. A. 2004, ApJ, 616, 998
  * Wilgenbus et al. (2000) Wilgenbus, D., Cabrit, S., Pineau des Forêts, G., & Flower, D. R. 2000, A&A, 356, 1010
  * Woitke et al. (2009) Woitke, P. et al. 2009, in Cool Stars, Stellar Systems, and the Sun: Proc. of the 15th Cambridge Workshop, AIP Conf. Proc., 1094, 225
  * Zinnecker et al. (1998) Zinnecker, H., McCaughrean, M. J., & Rayner, J. T. 1998, Nature, 394, 862

Figure 1: H2 $v=1-0$ $S(1)$ line spectra. Spectra acquired on 2007 Jun 24 and
2008 Jan 24 (first epoch) appear in the left panel, and spectra acquired on
2007 Jun 25 and 2008 Jan 25 (second epoch) appear in the right panel. These
spectra were extracted only in the region containing the mostly point source
continuum emission; any extended H2 emission is not included.

Figure 2: Histogram of mean H2 $v=1-0$ $S(1)$ line velocities minus
photospheric velocities for all observations.

Figure 3: Histogram of H2 $v=1-0$ $S(1)$ line FWHM velocities. The
instrumental line width of 17 km s-1 has been subtracted in quadrature from
each value before binning.

Figure 4: Histogram of H2 $v=1-0$ $S(1)$ line luminosities.

Figure 5: Histogram of H2 $v=1-0/2-1$ $S(1)$ line ratios. Values of 1.9, 7.7,
and 17 are indicated, nominally correspond to UV, shock, and X-ray secondary
electron impact excitation respectively in the presence of no dust (Gredel &
Dalgarno, 1995). These values are the inverse of those presented in column 8
of Table 2.

Table 1: Journal of Observations  Object | Region |  $\alpha$(J2000) |  $\delta$(J2000) | UT Date | Int. Time | Slit PA
---|---|---|---|---|---|---
|  | (hh mm ss.s) | ($\arcdeg$ $\arcmin$ $\arcsec$) |  | (minutes) | ($\arcdeg$E of N)
03260+3111B | Per | 03 29 07.7 | 31 21 58 | 2008 Jan 24 | 20.0 | 103
|  |  |  | 2009 Jan 25 | 20.0 | 149
03260+3111A | Per | 03 29 10.7 | 31 21 59 | 2008 Jan 24 | 20.0 | 118
|  |  |  | 2008 Jan 25 | 12.0 | 180
04108+2803B | Tau-Aur | 04 13 54.9 | 28 11 31 | 2008 Jan 24 | 20.0 | -123
|  |  |  | 2008 Jan 25 | 25.0 | 121
04158+2805 | Tau-Aur | 04 18 58.2 | 28 12 24 | 2008 Jan 24 | 20.0 | 73
|  |  |  | 2008 Jan 25 | 20.0 | 63
04181+2654AB | Tau-Aur | 04 21 11.5 | 27 01 09 | 2008 Jan 24 | 10.0 | 48
|  |  |  | 2008 Jan 25 | 30.0 | 172
DG Tau | Tau-Aur | 04 27 04.8 | 26 06 17 | 2008 Jan 25 | 10.0 | 57
04264+2433 | Tau-Aur | 04 29 30.0 | 24 39 56 | 2009 Jan 24 | 12.0 | 57
|  |  |  | 2008 Jan 25 | 12.0 | -106
04295+2251 | Tau-Aur | 04 32 32.1 | 22 57 27 | 2009 Jan 24 | 40.0 | -106
|  |  |  | 2008 Jan 25 | 30.0 | 75
04361+2547 | Tau-Aur | 04 39 13.5 | 25 53 20 | 2008 Jan 24 | 12.0 | 68
|  |  |  | 2008 Jan 25 | 12.0 | -110
04365+2535 | Tau-Aur | 04 39 35.2 | 25 41 45 | 2008 Jan 24 | 44.0 | 97
|  |  |  | 2008 Jan 25 | 33.3 | 99
GSS 30 | Oph | 16 26 21.4 | -24 23 06 | 2007 Jun 25 | 14.0 | -124
GY 21 | Oph | 16 26 23.6 | -24 24 38 | 2007 Jun 24 | 20.0 | -56
|  |  |  | 2007 Jun 25 | 20.0 | -100
WL 12 | Oph | 16 26 44.1 | -24 34 48 | 2007 Jun 24 | 8.0 | -47
|  |  |  | 2007 Jun 25 | 10.0 | -114
WL 6 | Oph | 16 27 21.6 | -24 29 51 | 2007 Jun 24 | 20.0 | -72
|  |  |  | 2007 Jun 25 | 30.0 | -89
IRS 43 | Oph | 16 27 27.0 | -24 40 50 | 2007 Jun 24 | 45.0 | -78
|  |  |  | 2007 Jun 25 | 60.0 | -66
YLW 16A | Oph | 16 27 27.8 | -24 39 32 | 2007 Jun 24 | 12.0 | -85
|  |  |  | 2007 Jun 25 | 12.0 | -133
IRS 67 | Oph | 16 32 01.1 | -24 56 45 | 2007 Jun 24 | 60.0 | -59
SVS 2 | Ser | 18 29 56.8 | 01 14 46 | 2007 Jun 25 | 28.0 | -85

Table 2: Protostar H2 Line Analysis Source | UT Date | 1–0 S(0) | 1–0 S(1) | 2–1 S(1) | AvbbV magnitude extinction was computed using each objects 2MASS $JHK$ colors, an extinction law, and estimated intrinsic CTTS locus colors as explained in the text in §3.4 | 1–0 S(1) | 2–1/1–0 | FWHMcc The mean FWHM velocity of the 1–0 S(1) H2 line, where the intrinsic instrumental line width of 17 km s-1 has been removed in quadrature. | V(H2 \- *)dd The radial velocity of the stellar photosphere subtracted from the mean radial velocity of the 1–0 S(0), 1–0 S(1), and 2–1 S(1) H2 emission lines (uncertain lines not used). No data indicate that the star lacked either H2 or photospheric lines. | H2 1–0 S(1)
---|---|---|---|---|---|---|---|---|---|---
|  | EW (Å)aaPositive equivalent widths indicate emission. Uncertainties are $\sim 0.02\AA$ for 2007 data and $\sim 0.05\AA$ for 2008 data. | EW (Å)aaPositive equivalent widths indicate emission. Uncertainties are $\sim 0.02\AA$ for 2007 data and $\sim 0.05\AA$ for 2008 data. | EW (Å)aaPositive equivalent widths indicate emission. Uncertainties are $\sim 0.02\AA$ for 2007 data and $\sim 0.05\AA$ for 2008 data. | (mag) | Log L(W) | S(1) | (km s-1) | (km s-1) | extent (″)
03260+3111B | 2008 Jan 24ee Spectrum displays $\Delta v=2-0$ CO emission. | 0.50 | 2.12 | 0.49 | 3.5 | 22.0 | 0.25 | 43 | … | $\gtrsim 10$
| 2008 Jan 25 | 0.51 | 2.69 | 0.39 | 3.5 | 22.1 | 0.14 | 44 | … | $\lesssim 1$
03260+3111AffExtended H2 emission is spatially displaced from the star. No H2 emission is coincident with the stellar continuum (as also observed by Doppmann et al. 2005). | 2008 Jan 24 | $<0.1$ | $<0.1$ | $<0.1$ | … | … | … | … | … | $\sim 7$
| 2008 Jan 25 | $<0.1$ | $<0.1$ | $<0.1$ | … | … | … | … | … | $\gtrsim 10$
04108+2803B | 2008 Jan 24 | 0.92 | 4.29 | 0.39 | 19 | 22.0 | 0.08 | 18 | -1 | $\sim 1$
| 2008 Jan 25 | 0.69 | 3.43 | 0.29 | 19 | 21.9 | 0.07 | 18 | -1 | $\sim 1$
04158+2805 | 2008 Jan 24 | 0.49 | 2.57 | 0.26: | 3.1 | 21.1 | 0.08: | 20 | -4 | $<1$
| 2008 Jan 25 | 0.35 | 2.70 | 0.22: | 3.1 | 21.2 | 0.06: | 20 | -4 | $\sim 2$
04181+2654AB | 2008 Jan 24 | 0.23 | 1.46 | 0.04:: | 23 | 22.0 | … | 24 | -1 | $\sim 1$
| 2008 Jan 25 | 0.39 | 1.37 | 0.00:: | 23 | 21.9 | … | 25 | 8 | $<1$
DG Tau | 2008 Jan 25 | 0.14 | 0.54 | 0.06:: | 0 | 22.0 | 0.08:: | 14 | -9 | $<1$
04264+2433 | 2008 Jan 24 | 1.56 | 5.78 | 0.53 | 6.1 | 21.7 | 0.06 | 10 | 4 | $\sim 2$
| 2008 Jan 25 | 1.44 | 6.09 | 0.52 | 6.1 | 21.7 | 0.06 | 13 | -4 | 6
04295+2251 | 2008 Jan 24 | 0.59 | 2.42 | 0.14: | 17 | 22.0 | 0.04: | 16 | … | $\sim 2$
| 2008 Jan 25 | 0.59 | 2.43 | 0.16: | 17 | 22.0 | 0.05: | 16 | … | $\sim 2$
04361+2547 | 2008 Jan 24 | 1.06 | 4.11 | 0.45 | 22 | 22.2 | 0.08 | 14 | … | 0
| 2008 Jan 25 | 0.66 | 2.10 | 0.21: | 22 | 21.9 | 0.07: | 13 | … | 0
04365+2535 | 2008 Jan 24 | 0.57 | 2.58 | 0.18: | 19 | 21.9 | 0.06: | 19 | … | 0
| 2008 Jan 25 | 0.35 | 1.59 | 0.09: | 19 | 21.6 | 0.05: | 16 | … | 0
GSS 30 | 2007 Jun 25 | 0.53 | 2.35 | 0.20 | 20 | 22.6 | 0.07 | 25 | -11 | 2
GY 21ggThe H2 emission slightly spatially displaced from the stellar continuum has velocity FWHM $\Delta v\sim 60$ km s-1, about twice that of the H2 emission spatially coincident with the stellar continuum. | 2007 Jun 24 | 0.37 | 1.80 | 0.20 | 16 | 21.8 | 0.07 | 28 | -11 | $\sim 1$
| 2007 Jun 25 | 0.52 | 2.21 | 0.20 | 16 | 21.9 | 0.07 | 29 | -10: | $\sim 1$
WL 12 | 2007 Jun 24 | 1.43 | 12.4 | 0.79 | 18 | 22.1 | 0.06 | 32 | -2 | $\sim 2$
| 2007 Jun 25 | 2.06 | 11.0 | 0.73 | 18 | 22.1 | 0.06 | 29 | -12 | $\sim 1$
WL 6 | 2007 Jun 24 | 0.06: | 0.34 | 0.06: | 37 | 21.6 | 0.15: | 15 | 0 | 0
| 2007 Jun 25 | 0.09 | 0.37 | 0.03: | 37 | 21.6 | 0.07: | 16 | -3 | 0
IRS 43hhExtended H2 emission shows velocity structure with a maximum FWHM $\Delta v\sim 40$ km s-1, about twice that of the H2 emission spatially coincident with the stellar continuum. | 2007 Jun 24 | 0.09 | 0.28 | 0.02:: | 33 | 21.8 | 0.05:: | 9 | 0 | $\sim 3$
| 2007 Jun 25 | 0.17 | 0.64 | 0.04: | 33 | 22.1 | 0.05:: | 22 | -4 | $\sim 3$
YLW 16A | 2007 Jun 24 | 0.32 | 1.41 | 0.20 | 17 | 21.6 | 0.12 | 25 | 0 | $\sim 2$
| 2007 Jun 25 | 0.31 | 1.25 | 0.13 | 17 | 21.6 | 0.09 | 28 | -4 | $\sim 3$
IRS 67 | 2007 Jun 24 | 0.13 | 0.37 | 0.03:: | 22 | 21.1 | 0.07:: | 15 | 0 | $\sim 1$
SVS 2iiThe extended H2 emission is separated from the stellar continuum and its line emission by $\sim 2\arcsec$. The extended emission has FWHM $\Delta v\sim 24$ km s-1, about twice that of the H2 emission spatially coincident with the stellar continuum. H2 line strength ratio was not calculated due to the very low value and high uncertainty ($\sim 50$%) of the 2–1 S(1) line measurement. | 2007 Jun 25 | 0.09 | 0.39 | 0.01::: | 0 | 21.5 | … | 13 | 19 | $\sim 3$