# A rotating accretion disk around MWC 297, a young B1.5 Ve star

Göran Sandell  Institute for Astronomy, University of Hawai‘i at Manoa, Hilo,
640 N. Aohoku Place, Hilo, HI 96720, USA<EMAIL_ADDRESS>William Vacca
NSF s NOIRLab, 950 N. Cherry Avenue, Tucson, AZ 85719, USA
<EMAIL_ADDRESS>
###### Abstract

High resolution spectra with iSHELL on IRTF in the K and M band of the young,
heavily accreting B 1.5e star MWC 297 show numerous double-peaked CO lines.
These CO lines originate in an inclined gaseous disk in Keplerian rotation.
MWC 297 is the only early B star known to show a Keplerian disk in CO.
Analysis of the spectra show that 12CO 1 – 0 is optically thick for the low
excitation lines. Even the 13CO 1 – 0 and 12CO 2 – 1 have somewhat optically
thick lines at low J levels. We find that the CO emission in the disk can be
fitted with CO being in a narrow ring at a radius of 12 AU, with a temperature
of 1500 K, and a CO column density of 1.6 1018 cm-2. This model underestimates
the line strength of high J lines, indicating that they are excited by
fluorescence. The CO overtone lines have a similar temperature, The 13CO lines
are much brighter than expected from interstellar isotope ratios. The 13CO
lines are wider than the 12CO ones suggesting different excitation conditions.
The same is true for 12CO – 1. We see strong absorption in 12CO and 13CO 1 – 0
at low J levels, which is due to two two cold foreground clouds. These clouds,
one with a temperature of 8.3 K and a column density of 6.7 1017 cm-2, and the
other one colder and with lower column density, can fully account for the
observed extinction toward MWC 297.

Star formation (1569); Star forming regions (1565); Massive stars (732);
Circumstellar disks (235); Infrared astronomy (786)

††thanks: Visiting Astronomer at the Infrared Telescope Facility, which is
operated by the University of Hawaii under contract
80HQTR19D0030 with the National Aeronautics and Space
Administration.

##  1 Introduction

Most young, low mass stars are surrounded by circumstellar disks. The same is
also true for intermediate mass stars up to a spectral type of $\sim$B8
(Williams & Cieza, 2011). Direct detections of disks around early B or O stars
are rare. This does not mean that high mass stars do not have accretion disks;
rather it indicates that disks are much more short-lived around high mass
stars so that by the time these stars become optically visible they have
already dispersed their accretion disks. Furthermore, high mass stars are
generally more distant than low mass stars, making it harder to spatially
resolve their accretion disks. Even O stars with stellar masses of 20 – 30 M⊙
appear to have disks in their heavy accretion phase (see e.g., Ilee et al.,
2013; Beltrán & de Wit, 2016; Zapata et al., 2019; Sandell et al., 2020;
Moscadelli et al., 2021). However, such stars are difficult to study because
they are always in clusters and deeply embedded.

Most of the evidence for disks around high mass stars comes from deeply
embedded young objects observed with interferometric means, mostly in the
millimeter/submillimeter regime with ALMA and NOEMA and other array
telescopes. There are also some detections at optical/infrared wavelengths
(Beltrán & de Wit, 2016). There has been some success in looking for Keplerian
rotation in accretion disks using the rovibrational CO lines at 2.3 $\mu$m and
at 4.5 - .2 $\mu$m. CO overtone emission ($\upsilon$ = 2 – 0 and $\upsilon$ =
3 – 1) was first detected in the BN object (mass $\sim$ 10 M⊙) (Scoville et
al., 1979), and in the early B-star MWC 349 A (Kraus et al., 2000). The latter
has a mass of $\sim$ 30 M⊙, although a recent paper by Kraus et al. (2020)
argue that it is a B[e] supergiant. There have been a number of detections of
bandhead emission toward massive young stellar objects (Ishii et al., 2001;
Blum et al., 2004; Bik & Thi, 2004; Bik, Kaper & Waters, 2006; Wheelwright et
al., 2010; Cooper et al., 2013; Ilee et al., 2013; Pomohaci et al., 2017). For
these objects the stellar mass is usually deduced from the observed bolometric
luminosity. These can be highly uncertain because it includes emission from
the accretion process and possible emission from nearby objects, resulting in
an overestimate of stellar mass, see, e.g., W 33 A, for which Pomohaci et al.
(2017) determine a mass 17.2 M⊙, while Navarete et al (2021) find a mass of 10
M⊙. Yet it is clear that some of these are indeed very young high mass stars.
The CO fundamental ($\upsilon$ = 1 - 0) rovibrational lines, located in the M
band, 4.5 – 5.2 $\mu$m, have been detected in low mass stars (Najita, Carr &
Mathieu, 2003), HAEBE stars, and transitional disks objects (Brittain et al.,
2003; Blake &Boogert, 2004; Salyk et al., 2009; van der Plas et al., 2015;
Doppmann et al., 2017; Banzatti et al., 2022), but not in high-mass stars (M
$>$ 8 M⊙), except for MWC 297, the subject of this study.

MWC 297 is a bright, nearby HAeBe star with a spectral type B1.5 Ve (Drew et
al., 1997) and a mass of $\sim$ 10 M⊙, located at a distance of 418 pc (Gaia
DR3, Riello et al., 2021). It has been the subject of several studies with
ESO’s Very Large Telescope Interferometer (VLTI) (Acke et al., 2008; Weigelt
et al., 2011; Lazareff et al., 2017; Kluska et al., 2020). Most of these
studies suggest that the star is surrounded by an almost face-on disk. The
star, however, drives an ionized outflow, with spatially separated outflow
lobes, and therefore cannot be face-on (Sandell, Weintraub & Hamidouche,
2011). Mid-infrared imaging with FORCAST on SOFIA (Vacca &Sandell, 2022) shows
that hot dust traces the outflow lobes at wavelengths $>$ $\mu$m. Simple
geometrical modeling the mid-infrared morphology constrain the disk
inclination rather well and give an inclination angle i = 55∘.

Here we present and discuss high resolution M band data obtained with the
iSHELL instrument at the IRTF. The spectrum reveals double-peaked CO emission
lines, which we have modeled with a rotating Keplerian disk to determine the
properties of the emission region. We also discuss and analyze the absorption
spectra from the cold foreground cloud, which are seen in low $\upsilon$ = 1 –
0 P and R transitions. We note that MWC 297 was included in a large survey of
planet forming disks observed with the same instrument and the same wavelength
range, and these data have recently been published (Banzatti et al., 2022).
While Banzatti et al. (2022) did classify the double peaked CO profiles in MWC
297 as originating in a Keplerian accretion disk, they provide very few
details. Here we present a more comprehensive analysis of of both CO overtone
emission at 2.3 $\mu$m as well as the rovibrational spectra in the M-band.

##  2 iSHELL observations and data reduction

Observations of MWC 297 were obtained at the NASA Infrared Telescope Facility
(IRTF) on Mauna Kea on 2020 Sep 30 (UT) and on 2022 April 26 (UT) with iSHELL,
the facility near-infrared high resolution cross-dispersed spectrograph
(Rayner et al., 2022). The M band observations used the M1 setting of iShell
both in 2020 and 2022. This mode yields spectra spanning the wavelength range
.5 – 5.2 $\mu$m over 16 spectral orders. The observations were acquired in
“pair mode”, in which the object was observed at two separate positions along
the 15′′-long slit. The slit width was set to 0$\farcs$375, which yields a
nominal resolving power of 88,000 for the spectra. (At the distance of MWC 297
of 418 pc, the iShell 0$\farcs$375 slit spans $\sim 160$ AU.) Twenty
individual exposure of MWC 297, each lasting 28 s, were obtained using the M1
setting of iShell. In 2020 the slit was set to the parallactic angle (46.0∘)
during the observations. In 2022 the observations were done using two slit
positions, 197∘ and 17∘, which were coadded in the post reduction. The
observations both times were acquired in “pair mode”, in which the object was
observed at two separate positions along the 15′′-long slit. The slit width
was set to 0$\farcs$375, which yields a nominal resolving power of 88,000 for
the spectra. At the distance of MWC 297, 418 pc, the iShell 0$\farcs$375 slit
spans $\sim 160$ AU. Long observations of an A0 V star, used as a “telluric
standard” to correct for absorption due to the Earth’s atmosphere and to flux
calibrate the target spectra, were obtained immediately prior to the
observations of MWC 297. In 2022 we also observed MWC 297 in the K band. These
observations were done in the K3 setting, providing a spectral coverage from
2.256 to 2.55 $\mu$m. A set of internal flat fields and arc frames were
obtained immediately after the observations of MWC 297 for flat fielding and
wavelength calibration purposes.

The data were reduced using Spextool (Cushing, Vacca & Rayner, 2004), the IDL-
based package developed for the reduction of SpeX and iShell data. The
Spextool package performs non-linearity corrections, flat fielding, image pair
subtraction, aperture definition, optimal extraction, and wavelength
calibration. The sets of spectra resulting from the individual exposures were
median combined and then corrected for telluric absorption and flux calibrated
using the extracted A0 V telluric standard spectra and the technique and
software described by Vacca, Cushing & Rayner (2003). The spectra from the
individual orders were then spliced together by matching the flux levels in
the overlapping wavelength regions, and regions of poor atmospheric
transmission were removed. However, for the September 2020 the final M
spectrum of the A0V star had relatively poor S/N (on the order of 30), and
this limited the final S/N of the spectrum of MWC 297, which is considerably
brighter in the M band. (The reduced spectrum of MWC 297 has a S/N of several
hundred.) Therefore, an alternate telluric correction was derived using the
xtellcor_model software developed by Boogert (2022) and then applied to the
reduced spectrum of MWC 297. Although this was effective at correcting the
telluric absorption, the resulting spectrum was still uncalibrated. Therefore,
we fit a polynomial to the continuum of both the A0V-corrected spectrum and
the model-corrected spectrum, and multiplied the model-corrected spectrum by
the wavelength-dependent ratio of the polynomials. This yielded a telluric-
corrected and flux-calibrated spectrum of MWC 297. The S/N varies across the
spectrum but is on the order of several hundred across the entire M1
wavelength range. For the 2022 data we used a brighter A0 V star, which could
directly be used for flux calibration. Direct comparison of the two calibrated
data sets show that the calibration agrees within 10% for both data sets and
we have therefore added together the two data sets, after correcting both
spectra for Doppler shifts due to the relative motion of the earth and the
source, and extracted flux densities for all lines not affected by telluric
absorption. In 2020 the Doppler shift for MWC 297 was -11.5 km s-1, which
severely affected CO lines with J values $\lesssim$ 15, while it was +41.2 km
s-1 in 2022, which largely shifted all CO lines away from telluric absorption
features.

##  3 A short overview of what has been learned from studies of rovibrational
CO

Although CO vibrational overtone emission was first detected in the inner
disks of both high and low mass objects, but it is not as ubiquitous as the CO
fundamental ($\Delta\upsilon$ = 1) transitions of CO at 4.6 $\mu$m. The
detection rate of overtone emission in unbiassed surveys is about 25%. In part
this is due to the relatively low A-values, which require high column
densities of hot gas for the lines to be detectable, but other factors play a
role as well. Modeling by Ilee et al. (2018) suggest that moderate accretion
rates produce the most prominent and hence detectable CO first overtone
emission. If the accretion rate is too high ($\gtrsim$ 10-4 M⊙ yr-1), it
results in large dust sublimation radii, a larger contribution of hot dust to
the K band continuum and therefore a lower CO to continuum ratio. Low
accretion rates ($\lesssim$ 10-6 M⊙ yr-1) however, result in smaller dust
sublimation radii, a smaller CO emitting area and therefore a lower CO to
continuum ratio. Although CO overtone emission is detectable in MWC 297, which
has an accretion rate of 3 10-4 M⊙ yr-1, the first overtone band is barely
visible in high resolution spectra, suggesting that the gas is relatively
cold, or about 1000 K. The higher vibrational levels are filled with double
peaked spectra. This will be discussed later in the paper.

The rovibrational CO line emission at 4.5 – 5.2 $\mu$m is as already mentioned
much more ubiquitous and seen in a variety of objects, ranging from Transition
Objects to Classical T Tauri and HAEBE stars, but contrary to CO overtone
emission it is not seen in high-mass stars except for MWC 297 (Banzatti et
al., 2022, and this paper) and perhaps in BN (Scoville et al., 1983). Banzatti
et al. (2022) gives an up to date list of all the high resolution K and M-band
CO surveys done over the last 20 years. Therefore there is no reason for us to
repeat it here. Instead we try to surmmarize the major findings resulting form
these studies and complementary theoretical modeling.

The rovibrational CO emission at 4.7 $\mu$m is very common in protoplanetary
disks and arises in the inner 0.1 - 2 AU region in disks around low mass stars
Najita et al. (2007) and up to 30 AU or more in HAEBE stars (van der Plas et
al., 2015; Banzatti et al., 2018, 2022). The line shapes vary (Bast et al.,
2011; Brown et al., 2013; Banzatti & Pontoppidan, 2015; Banzatti et al.,
2022). They can be narrow and single peaked, narrow with a broad pedestal,
broad but single peaked, or pure double peaked Keplerian profiles. Some stars
also show the CO in absorption, which in most cases originates from colder
foreground gas. Double peaked profiles, which are expected to be a true disk
signature, are surprisingly uncommon (Bast et al., 2011), although in the
large survey by Banzatti et al. (2022) the detection rate of double peaked
lines was quite common, $\sim$ 50%.

Banzatti et al. (2022) define the line profiles in two broad groups depending
on line shape: triangular and double peaked lines. In their classification the
single peaked narrow lines with a broad pedestal and the broad, but single
peaked lines all fall into this category. These lines are often somewhat
asymmetric with the blue-shifted side being stronger Herczeg et al. (2011);
Brown et al. (2013). (Banzatti et al., 2022) define a line shape parameter S
as S = FW10%/FW75%, where FW stands for Full Width at 10 and 75%,
respectively. Lines with a shape value $\gtrsim$ 2.5 are triangular, and
$\lesssim$ 2.5 they are generally double peaked Keplerian or narrow single
peaked lines. The latter are believed to represent Keplerian disk emission
seen in face on disks. The single peaked narrow lines with a broad pedestal,
i.e. triangular shaped lines cannot be explained by Keplerian motion. The
broad pedestal may still be Keplerian, but the narrow component appears to
originate in a slow mowing disk wind Bast et al. (2011); Brown et al. (2013).
Pontoppidan et al. (2011) found from spectroastrometry of narrow line with a
broad pedestal that the narrow line emission is too compact and asymmetric to
originate in a Keplerian disk. The broad single peaked lines are also believed
to be a combination of a Keplerian disk rotation and a disk wind, but the
disks are seen with higher inclination angles, producing broader lines (Brown
et al., 2013).

Line shape alone does not really discriminate where the CO lines are formed,
whether they originate from the disk or the outflow or both. Spectroastrometry
(Pontoppidan et al., 2011; Brittain, Carr & Najita, 2018) can achieve
subarcsecond resolution and hence put spatial constraints on the emitting
region, which has been used to discriminate between Keplerian and wind models.
Another powerful constraint is the strength of higher vibrational levels,
which have been detected up to $\upsilon$ = 6 – 5 in HAEBE stars (Brittain et
al., 2007; Jensen et al., 2021; Banzatti et al., 2022). For double-peaked
lines, the work by Banzatti et al. (2018, 2022); Bosman et al. (2019) suggests
that one can identify three emitting regions that show a dependence on the
$\upsilon$2/$\upsilon$1 ratio depending on the location of their emitting
radii, $R_{\rm CO}$, compared to their dust sublimation radius, $R_{\rm
subl}$. Regions a and 1b both have high vibrational ratio
$\upsilon$2/$\upsilon$1 = 0.2 – .0, but come from very different regions in
the disk, although both are largely dust free. Region 1a is inside $R_{\rm
subl}$, while 1b is inside a dust cavity observed by millimeter
interferometry. Region 1b therefore has lower rotational temperatures and
narrower line widths and it appears that UV fluorescence largely excites the
higher vibrational levels in this region (Bosman et al., 2019). In region 2
the vibrational ratio $\upsilon$2/$\upsilon$1 = 0.01 – 0.2, and the CO is
emitted from regions outside the dust rim.

##  4 Analysis and Results

The iSHELL spectrum exhibits a large number of double-peaked emission lines
from rovibrational transitions of CO. The $\upsilon$ = 1 – 0 CO lines up to
rotational level P(47) and higher vibrational levels up to $\upsilon$ = 5 – 4
are clearly seen. Emission from 13CO is also seen, in both the $\upsilon$ = –
0, and $\upsilon$ = 2 - 1 transitions. The profiles of the lowest rotational
transitions exhibit strong narrow absorption lines, due to absorption from a
cold foreground cloud, which cuts out a substantial fraction of the emission.
Three H lines (Pf $\beta$, Hu $\epsilon$, and Hu $\delta$) are also seen in
emission. The K band iSHELL spectrum also shows a large number of double
peaked lines, all from CO overtone bandhead emission, see Figure 1.

It is clear that the double peaked 12CO profiles must originate in a rotating
accretion disk. To enable a better comparison between 12CO and 13CO lines we
generated median line profiles of the 12CO $\upsilon$ = – 0, 13CO 1 – 0 and
also for the 12CO 2 – 1 and 12CO 2 – 0 emission lines from the spectrum after
subtracting a low-order polynomial continuum. Since we see no difference
between low and high J lines we created median lines profiles of all
unaffected by absorption or blending. For the 12CO 1 – 0 we used R(14), R(9),
R(6), P(8), P(9), P(12), P(14), P(26), P(34), (P37), P(38), P(45), and P(47).
For 13CO 1 – 0 we used R(29), R(24), R(22), R(21), R(19), R(16), R(15), R(13),
R(12), R(10), R(9), R(3), P(3), P(7), P(9), P(11), P(12), P(16), P(17), and
P(21). The 12CO 2 – 1 transitions were R(27), R(21), R(20), R(17), R(14),
R(13), R(12), R(9), R(1), P(1), P(2), P(6), P(8), P(15), P(21), P(25), P(26),
P(28), P(32), and P(33), while the transitions for 12CO 2 – 0 were R(6), R(7),
R(11), R(12), R(14), R(17), R(18), R(20), R(21), R(22), R(23), and R(24).
These were re-interpolated and centered at the theoretical line center, see
Figure 2. From this plot it is easy to that the 13CO 1 – 0 lines are broader
and the flank of the 13CO lines also appear somewhat steeper. The 12CO 2 – 1
lines are also broader than the 1 – 0 lines. The overtone 12CO 2 – 0 lines are
even broader than any of the rovibrational lines. The separation between the
line peaks are also wider for both 13CO 1 – 0 and 12CO 2 – 1 compared to 12CO
1 – 0 lines indicating that these lines are inward of the disk relative to the
location of the 12CO 1 – 0 emitting gas. In Table 1 we give the measured FWHMs
and peak separations for all the lines shown in Fig. 2.

At 10% level the line widths are similar or $\sim$ 65 km s-1, suggesting that
the inner radius is similar for all of them. The FW10% maybe slightly larger
for the 12CO 2 – 0 lines. We measure a linewidth of 68 km s-1. It is not clear
whether this difference is real or simply due to measurement inaccuracy,
especially since Banzatti et al. (2022) quote a FW10% of 68.9 km s-1. They
also quote a slightly larger FWHM of 46.8 km s-1 for 12CO 1 – 0 than what we
find, 45.3 km s-1. We determine a Vlsr = +6.7 $\pm$ 2.2 km s-1for the emission
lines. This is a weighted mean of the September 2020 and April 2022 data sets.
Banzatti et al. (2022) find a center velocity of + 7 km s-1 (converted to
LSR). which agrees very well with what we find.

Table 1: Line profile parameters for median averaged emission lines. The FWHM values have measurement errors of $\sim$ 0.5 km s-1, the errors for the peak separations are $\sim$ 1 km s-1. Profile | FWHM | Peak separation
---|---|---
| [km s-1] | [km s-1]
12CO $\upsilon$ = 1–0 | 45.3 | 29
12CO $\upsilon$ = 2–1 | 48.2 | 32
13CO $\upsilon$ = 1–0 | 48.2 | 32
12CO $\upsilon$ = 2–0 | 54.7 | 37
Figure 1: The CO overtone bands observed with iSHELL in K band showing a large
number of double peaks lines. The 2 – 0 bandhead at 2.29 $\mu$m is very weak.
Figure 2: Median line profiles of 12CO $\upsilon$ = 1 – 0 and 2 – 1 and 13CO –
0 and 12CO 2 – 0 emission lines. The dashed blue line is the 12CO 2 – 1
profile scaled by 2.25, the dashed red line is the 13CO 1 – 0 profile scaled
by 3, the dashed orange is the scaled 12CO 2 – 0. The scaled profiles
demonstrate the differences in widths compared to the 12CO –0 profile.

###  4.1 Modeling the emission lines with a rotating Keplerian disk.

Figure 3: Rotational diagram with the observed data points plotted with large
dots. The line bands are identified by their colors in the bottom left part of
the plot. For almost all of the data, the error bars are smaller than the dot
size.  Figure 4:  The ring model for a rotational temperature of 1500 K and a
density of $1.6\times 10^{18}$ cm-2 plotted with solid lines on the data
points plotted in Figure 3. The model matches the 1 – 0 data quite well except
at the highest energies where it underestimates the line strength. It also
underestimate the line strength at the higher vibrational states, suggesting
that the excitation of these lines are dominated by fluorescence. The model
overestimates the – 1 data, most likely because its emitting area is smaller
than that of 1 – 0. The model calculations show that the 13CO abundance needs
to be increased by a factor of 6 - 7 compared to normal interstellar
abundances.

Figure 5: The ring model for T = 1500 K and a log density of 18.2 cm-2 in red
overlaid on the observed M band spectrum. 13CO needs to be scaled by 6 to
match the line profiles. The line transitions are marked and the rovibrational
bands they belong to are indicated at the top for 12CO and at the bottom shows
13CO 1 – 0. Regions of strong telluric absorption have been blanked out.

Inspection of the K and M band spectra yields conflicting information
regarding the CO emitting gas. The relative strengths of the 12CO lines across
the M band suggest that the gas is optically thick. In addition, the 13CO
emission lines are much stronger relative to those of 12CO than expected given
the normal isotopic ratio and optically thin gas. However, the clearly double-
peaked profiles suggest the emission cannot be completely optically thick.

In order to determine the properties of the CO emitting gas, we constructed
models of the emission regions under the assumption that the gas was in
Keplerian orbit around a 10 $M_{\odot}$ central source. We considered two
configurations: a uniform (i.e., constant temperature and surface density)
ring and a thin disk whose temperature and surface density varied with radial
distance. We adopted an inclination of 55∘ derived by Vacca &Sandell (2022).
The models assume the gas is in LTE and are similar to those constructed by
Kraus et al. (2000) and others. We generated a synthetic spectrum of the M
band emission for each value of temperature between 900 and 2000 K in steps of
100 K and log CO surface density between 15.0 and 22.0 in steps of 0.2. We
assumed a source distance of 418 pc and the standard isotopic ratio for
12CO/13CO of 77 (Wilson & Rood, 1994). For the ring models we adopted a radius
of 12 AU in order to match the mean width of the emission lines, and a ring
width of 1 AU. For the disk models we adopted a power law index of 0.4 for the
temperature and 1.5 for the surface density and radial steps of 0.5 AU between
an inner radius of 5 AU and an outer radius of 25 AU.

We then reddened each model by $A_{V}=8.1$ mag, convolved it with a Gaussian
whose width is given by the instrumental resolution ($R\sim 88,000$ in the M
band), and sampled it at the observed wavelength sampling of the observations.
A $\chi^{2}$ comparison between the observed 12CO line spectra and the model
spectrum was then computed. For the ring model we computed the best fit
scaling of the model to the spectroscopic data, which corresponds to the best
fit width of the ring. The emission line strengths were also compared with the
values measured from the observed spectra on the standard rotational diagram
(Figure 4).

Somewhat surprisingly, we found that no disk models with our assumed
parameters were able to provide good matches to the observed spectra. A thin
ring model with a temperature of $1500$ K and a surface density of $\log
N_{CO}=18.2$ provided a reasonable, although not perfect, fit to the data. It
was immediately clear, however, that the 13CO emission was far weaker in the
model than in the actual data by a factor of 6 - 7. This was also reflected in
an offset between the points for the 13CO emission line strengths from the
model and that of the data on the rotational diagram, although the slope of
the points for the 13CO emission line strengths from the model matched that of
the data (which indicates that the derived temperature was approximately
correct). This result, in addition to the differences seen in the line widths
of the 12CO and 13CO emission profiles, suggests that the 12CO and 13CO
emission arises from different parts of the disk. Therefore, we generated a
pure 13CO emission model in which the 13CO gas temperature was the same as
that for the 12CO gas, but the abundance was enhanced by a factor of 6 , which
gave a good fit to the 13CO 1 – 0 spectra. Figure 5 shows the observed spectra
overplotted with the best fit model without any additional scaling for 13CO.
The corresponding line strengths are shown on the rotation diagram in Figure
4.

Despite the excellent agreement between the model and the line strengths seen
in Figure 5, it is clear that the best-fit model is somewhat unsatisfactory. A
thin ring of CO emission seems physically unrealistic. Furthermore, although
the model provides a reasonable match to the overall line strengths, it does
not provide a good match to several prominent lines in the spectrum. Many of
the 12CO 2–1 lines in the model are much stronger than in the data. In
addition, the line profiles of many of the low lying 12CO – 0 transitions are
not well reproduced by the model.

We have not done a detailed analysis of the $K$-band overtone spectrum, but
the – 0 bandhead emission is very weak and the 3 – 1 bandhead is almost
missing, indicating that the CO gas is quite cold. The best fit model from the
analysis of the rovibrational CO lines provides an adequate fit to the data.

###  4.2 Absorption lines

The rovibrational fundamental CO lines in the low P and R transitions show
narrow absorption lines cutting out most of the emission from the accretion
disk in the center of the emission lines in the $M$ band spectra. Absorption
is also seen in the lowest R and P transitions of 13CO in both data sets. This
suggest that the emission is being absorbed by a cold foreground cloud. This
cold foreground cloud is likely to contribute to most of the extinction toward
MWC 297, which is estimated to be $\sim$ 8.1 mag (Vacca &Sandell, 2022).

The lowest transitions of the CO absorption lines in the 2020 data set are
strongly affected by telluric absorption, which makes it difficult to estimate
the true extent of the absorption. Therefore, our analysis focussed on the
April spectra, for which the relative velocity between the intrinsic
absorption and the telluric absorption is larger. In order to extract the
absorption lines we generated a mean profile for the $\upsilon$ = 1 – 0 12CO
emission lines and then scaled that to match the observed profiles of the low
J transitions that have the foreground absorption (see Figure 6). Examination
of the CO absorption lines indicates that the lowest transitions reach the
zero intensity level and appear much broader than the higher transitions,
suggesting that they are optically thick. The higher level J transitions agree
well in velocity with the 13CO absorption lines, while there is a clear shift
in velocity at the lowest levels. We illustrate this in Figure 7, which shows
three CO $\upsilon$ 1 – 0 lines and three of the 13CO lines in the source
spectrum, which has been shifted to 0 km/s. It can be seen from this figure
that the weaker absorption lines are blue-shifted relative to the disk
emission. and that the lines are asymmetric, with a red tail (Figure 7). For
the optically thin 12CO and 13CO lines we find a velocity of -0.6 $\pm$ 0.5 km
s-1. With increasing optical depth blue-shift of the lines decreases and the
red tail increases. At the lowest J levels the CO is very optically thick, the
profiles look flat topped and the center of the line is shifted to the red.
This suggests that there are two absorption components to these lines, with
the red shifted component contributing substantially to the width of the
lowest and most optically thick transitions. The radial velocity of MWC 297,
6.7 $\pm$ 2.2 km s-1, which we determined for the center of the emission lines
is in good agreement with the velocity of the molecular cloud in which it was
born. Sandell et al. (2023) determine a velocity $\sim$ 7.2 km s-1 of the
molecular cloud from 13CO(6–5) observations with APEX. This agrees well with
13CO(3-2) and C${}^{18}O$(3–2) observations. These lines are unaffected by the
cold foreground cloud, while all the emission in the low J 12CO emission is
completely absorbed (see also Manoj et al., 2007). Only emission from CO(4–3)
and higher rotational levels are unaffected by the cold foreground cloud

Figure 6: An example illustrating the extraction of absorption line profiles.
The observed emission line, in this case the 12CO $\upsilon$ = 1 – 0 R(4) is
plotted in black, the scaled template profile is plotted in red and the
resulting absorption profile is green. The baseline around the absorption
profile (green) is quite flat and centered at 0.  Figure 7:  12CO and 13CO
$\upsilon$ = 1 – 0 absorption lines. The 12CO profiles are plotted with solid
lines and 13CO with dashed lines. The lines are identified in the bottom left
corner of the image. The velocity of the emission lines, centered at 0 km s-1,
is marked with a gray vertical line. The optically thick 12CO $\upsilon$ = 1 –
0 R(0) line appears flat topped, much broader, and shifted in velocity
relative to 12CO 1 – 0 R(4) and 13CO 1 – 0 R(2), which are optically thin.
This is because there is a second, colder red-shifted cloud component, which
is optically thick in the 12CO 1 – 0 R(0) line. It is also seen in the 13CO 1
– 0 line as a strong red-shifted wing. The 0 km s-1 velocity corresponds to a
Vlsr = 6.7 km s-1.

To determine the temperature and column density of the absorbing gas we
performed a curve of growth analysis similar to that of Goto et al. (2003). We
solved for column density, temperature, and the line width using an iterative
least-squares fitting routine. However, the low excitation CO lines are very
optically thick and also include the red-shifted component, which we cannot
separate at the iShell resolution. We therefore first analyzed the weaker, and
presumably more optically thin, 13CO lines. Using all 13CO lines in the
analysis resulted in large densities (2 1018 cm-2) and very low temperatures
($\sim$4 K). While these parameters provided a reasonable match the observed
equivalent widths (EWs), the model did not reproduce the line profiles or the
depths. The fit is skewed to high densities in order to reproduce the observed
EWs, which are clearly dominated by the second red absorption component for
the highest optical depth lines. As can be seen in Figure 7, even the 13CO
lines exhibit a red wing, which is contributing to the EW. We therefore
restricted the EW analysis to the two lowest opacity lines, for which the red
wing appears to have the smallest contribution, and derived a column density
of $\sim$ 6.5 1017 cm-2 and a temperature of $\sim$ 8.5 K. We then combined
the two weakest CO lines with the two weakest 13CO lines in the fit and
obtained a solution that matched both the EWs and the line profiles. We kept
the velocity width equal to 3.5 km s-1for this fit, which is in aqreement with
our resolution (R $\sim$ 88,000). This resulted in a CO column density, N(CO)
= 6.7 $\pm$ 0.4 1017 cm-2, and a temperature T = 8.25 $\pm$ 0.10 K. No
substantial residuals are seen when we subtract this model from the spectrum
of the 4 lines. If we assume a CO abundance of 10-4, and that all the hydrogen
is molecular and adopt $N(H_{2})/A_{V}=0.94~{}10^{21}$ molecules cm-2 (Bohlin
et al., 1978), then this absorption component indicates an extinction of 6.3
mag. This, however, is a lower limit to the total foreground extinction since
the red-shifted absorption component is completely opaque at the lowest CO
transitions and also detected in 13CO. It is likely that it is responsible for
an additional extinction of 1 - 2 magnitudes. Therefore the cold foreground CO
emission readily explains the observed foreground extinction of 8.1 mag.

Figure 8:  iSHELL spectrum showing the $\upsilon$ = 1 – 0 12CO lines in
absorption. The wavelength region is split into two for clarity. The
transitions R(3), R(2), R(1) and R(0) at the top, P(1), P(2), P(3) and P(4) on
the bottom, left to right. Regions of strong telluric absorption have been
blanked out. The fainter double peaked emission lines are from higher
rotational rotation transitions. The strong hydrogen lines Pf $\beta$ and Hu
$\epsilon$ are also seen. The model fit to the absorption lines is overlaid in
red.

##  5 Discussion

Generally half of all solar mass stars have lost their disks after 3 Myr with
an overall disk life time of 6 Myr (Haisch, Lada & Lada, 2001). Some stars
still have disks after 10 Myr (see e.g. Espaillat et al., 2008). However, disk
dispersion timescales for high-mass stars are much shorter. By the time O and
B stars become optically visible, they have long since dispersed their disks.
It is therefore not surprising that hardly any disks have been detected around
visible early B stars. The two previously known examples are the early B star
MWC 349 A, and LkH$\alpha$ 101\. The latter is a B1 Ve star that illuminates
an H II region. Infrared interferometry shows that it is surrounded by an
almost face-on disk with an inner hole and a dust radius of $\sim$ 30 AU
(Tuthill et al., 2002). MWC 297 is now a third example. It has long been
suspected to have an accretion disk, but the disk size and inclination have
been uncertain. MWC 297 is a much younger star than MWC 349 A or LkH$\alpha$
101, as it is still heavily accreting with an age in the range of – 2.8
$\times$ 105 yr (Fairlamb et al., 2015; Vioque et al., 2018).

The rovibrational CO lines as well as the CO overtone bands show clear double
peaked profiles confirming that we see inclined circumstellar disk around MWC
297, which is in Keplerian rotation. MWC 297 does not really fit in the
classification scheme by Meeus et al. (2001) since it is much younger than
typical HAeBe stars, but if one follows the Meeus classification it would be
in group I (Guzmán-Días et a., 2021). For this reason Acke & van den Ancker
(2004) added a group III for deeply embedded and accretion dominated HAEBE
stars. Neither does MWC 297 readily fit into the classification of CO spectra
by Banzatti et al. (2018, 2022); Bosman et al. (2019). It has double peaked
profiles and the ratio of vibrational levels of $\upsilon$2/$\upsilon$1 = 0.45
it in region 1, but since there is no cold dust in the disk (Vacca &Sandell,
2022), we cannot definitely say whether it belongs to region 1a or 1b.
However, since the UV-fluorescence dominates the excitation of the $\upsilon$
= 2 – 1 and 3 – 2 bands and even appears to dominate the high excitation
levels of CO 1 – 0, it should probably go into group 1b. We also note there
are only five HAeBe stars that have overtone bandhead CO emission (Banzatti et
al., 2022), yet the bandhead emission was easily detected in MWC 297. Usually
the overtone emission comes from hot, dense CO gas in the innermost part of
the disk (Ilee et al., 2013, 2014) with gas temperatures of $\sim$ 3000 K and
column densities of 1021 cm-2. The overtone emission in MWC 297 is much
colder, $\sim$1500 K, an order of magnitude less dense, and extending up to a
radius of 10 AU.

It was clear from the modeling of the CO 1 – 0 lines that the line strengths
cannot be fit with a single temperature model. The is very little change in
line strength between low and high excitation lines and they all have the
characteristic double peaked profile from a disk in Keplerian rotation. An
isothermal fit to all the rotational lines results in very high optical depths
for the lowest transition and makes the line profiles flat-topped, which is
not what is observed. It is clear that the lines can only be moderately
optically thick, since they all show a clear U-shaped profile. Neither can a
single temperature model explain the strong $\upsilon$ = 2 - 1 and 3 – 2
transitions, which for the 2 – 1 lines are almost half the line strengths of
the 1 – 0 lines. This was already noticed by Brittain et al. (2003), who
argued that the strong $\upsilon$ = 2 - 1 and 3 – 2 transitions in isolated
HAeBe star HD 141569 must be excited by UV fluorescence, see also Brittain et
al. (2007). Recent modeling of HD 141569, which is a well studied HAeBe star
with vibrational levels detected up to $\upsilon$ = 7 – 6 confirm that
fluorescence is needed to explain the observed rovibrational lines (Jensen et
al., 2021) We expect MWC 297 to have strong UV radiation, both from the fact
that it is an early B star and because it is heavily accreting. However, due
to the high foreground extinction, 8 .1 mag, the star is not visible in the
UV. Therefore we have not attempted to include fluorescence in ir model,
although it it is clear that UV fluorescence is needed to explain the emission
in the higher vibrational bands and even necessary for the high excitation 1 –
0 lines. The same may be true for other HAeBe disks as well.

Because MWC 297 is very bright in the infrared, it has been the target of many
interferometric studies. These studies show a well resolved inner disk in the
continuum with a size of 2 – 4 AU in continuum (Malbet et al., 2007; Kraus et
al., 2008; Acke et al., 2008; Weigelt et al., 2011; Hone et al., 2017; Kluska
et al., 2020) with the line emission (Br$\gamma$) being 2 – 3 times more
extended (Malbet et al., 2007; Kraus et al., 2008; Weigelt et al., 2011; Hone
et al., 2017). Acke et al. (2008) also observed MWC 297 in the mid-inrared.
Modeling of mid-infrared data required a two component model and they found
that a smaller inner disk with a size of 4 AU and a more extended disk with a
size of 17 AU (40 mas), This is rather similar to the size of disk we see in
CO Neither do these studies find any inner gap (Acke et al., 2008; Kluska et
al., 2020). What is also noteworthy is that the size of the continuum region
is well within the dust destruction radius, which for MWC 297 is $\sim$ 5
AU111Kluska et al. (2020) quotes a dust destruction radius of 8 AU, but that
is because they use a stellar luminosity 4 $\times$ 104 L⊙, which is far too
high for a B1.5 V star, see Vacca &Sandell (2022).. As pointed out by Vacca
&Sandell (2022) the Br$\gamma$ emission is extremely optically thick and it is
therefore appears that the continuum comes from an optically thick gaseous
region. From the observed line width of the CO lines at 10% level, $\sim$ 65.2
km s-1, we find the inner CO radius $\sim$ 5.6 AU, i.e. roughly at the same
radius as the dust destruction radius,

A normal isotope ratio of [12C/13C] of 77 greatly underestimates the strength
of the observed 13CO line intensities. We find that a relative 12CO to 13CO
abundance of $\sim$ 6 provides a good fit to the data. This does not mean that
the disk has anomalous isotope ratios, but rather that 13CO is not in LTE and
coexisting with 12CO. As we seen, the FWHM of 13CO 1 – 0 differs by more than
5% from that of 1 – 0, see Section 4. These kind of 12CO/13CO ratios are not
uncommon. van der Plas et al. (2015) found that the 12CO/13CO ratio varied
between and 7 for their sample of 12 HAeBe stars and that the rotational
temperatures of 13CO were on the average lower than those for 12CO. Banzatti
et al. (2022), who has the largest sample of rovibrational CO lines in HAeBe
stars, do not quote a 12CO/13CO ratio, but looking at the line ratios for
their sample of 17 HAeBe stars with double peaked line profiles shows that the
ratio ranges from 3 to 25, with a median of 14. This does not directly
translate to an abundance ratio [12CO/13CO], because for most disks 12CO 1 – 0
is more optically thick than 13CO 1 –0 , nor does it account for difference in
excitation between 12CO and 13CO.

We detect narrow absorption lines, FWHM $\sim$ 4 - 5 km s-1 from a cold
unrelated foreground cloud in both 12CO and 13CO. Radio observations show that
this foreground cloud is very extended and therefore more local, probably
somewhere around 100 to 200 pc from us. At the lowest J transitions the
absorption lines in 12CO are completely saturated. At higher J levels the
lines become more blue-shifted and agree in velocity with the 13CO absorption
lines. The lowest transitions of 13CO is also optically thick and show a red-
shifted line wing.wing. Detailed analysis shows that the red-shifted component
comes from a colder, but less dense cloud component. For the blue-shifted
cloud component, which is at a Vlsr $\sim$ 6.4 km s-1 we derive a temperature
of 8.3 $\pm$ 0.1 K and a CO column density, N(CO) = 6.7 $\times$ 1017 cm-2. If
we assume that all the CO gas is molecular with a typical abundance ratio of
10-4 it corresponds to an extinction of .3 magnitude assuming the standard
relation between molecular hydrogen and extinction. Since the redshifted cloud
component also contributes to the extinction, it is clear that foreground
cloud plus diffuse gas along the line of sight can fully explain the observed
high extinction, 8.1 mag, toward MWC 297. Banzatti et al. (2022) noticed that
the absorption line FWHM decreased as a function of J level from 10 km s-1 in
the J = 1 lines down to a minimum of 3.3 km s-1 in the $\upsilon$ = 1 – 0 J =
4 lines and that the line shape also changed, becoming less Gaussian and
showing effects of saturation at the lowest J levels. They also saw a similar
behavior in 13CO. This agrees well with what we have found, and we now now it
is due to two cold clouds in the foreground of MWC 297 with different
temperatures and densities, but both of which are completely optically thick
in 12CO at the lowest J levels.

##  6 Summary and Conclusions

We have shown that overtone bandhead and rovibrational CO lines trace a hot
gas disk in Keplerian rotation around the 10 M⊙ star MWC 297. Our modeling
shows that the emission cannot be explained by thermal excitation alone. It
therefore appears that the high excitation rotational lines are largely
excited by fluorescence, something which has been seen in other HAeBe disks as
well. Analysis of the spectra show that 12CO $\upsilon$ = 1–0 emission is
optically thick for the low excitation lines. Even the 13CO 1 – 0 and 12CO 2 –
1 have somewhat optically thick lines at low J levels. We find that a narrow
ring with a radius of 12 AU and an inclination of 55∘ provides an adequate fit
to data. For this model we derive a rotational temperature of 1500 K and a CO
column density of 1.6 $\times$ 1018 cm-2. This model underestimates the line
strength of high J lines, indicating that these lines are excited by
fluorescence. The CO overtone emission, which is only seen in a few HAeBe
disks, has a similar temperature. Our best-fit model is still somewhat
unsatisfactory. A thin ring of CO emission seems physically unrealistic.
Furthermore, although the model provides a reasonable match to the overall
line strengths, it does not provide a good match to several prominent lines in
the spectrum. Many of the 12CO 2–1 lines in the model are much stronger than
in the data. We also find that a normal isotope ratio of [12C/13C] of 77
greatly underestimates the 13CO line intensities, while a 12CO to 13CO
abundance of $\sim$ 13 provides a good fit to the data. Other HAeBe disks
typically show similarly enhanced 13CO line intensities, most likely because
excitation conditions in the disks differ between 12CO and 13CO.

The 12CO 1 – 0 spectrum shows a strong absorption feature in the center of the
emission profile, which is completely saturated at J = 1 and 2, and which is
also seen in 13CO 1 – 0. This absorption is due to cold foreground emission,
which is unrelated to MWC 297. Radio observations show that this foreground
cloud is very extended and therefore more local, probably somewhere around 100
to 200 pc from us. Detailed analysis show that this foreground absorption is
caused by two cold foreground clouds, partly overlapping in velocity. The more
opaque one is slightly blue-shifted relative to the center of the emission
lines has a temperature of 8.3 K and a column density of 6.7 $\times$ 1017
cm-2, which correspond to a visual extinction of 6.3 mag. With our velocity
resolution we could not separate the second red-shifted component well enough
to be able to good estimate the temperature and density, but it is apparent
that it is colder and has a lower column density. It is clear these foreground
clouds are responsible for the high extinction toward MWC 297.

We have shown that the young, heavily accreting B1.5 V star MWC 297 is still
surrounded by a molecular accretion disk in Keplerian rotation. It is the only
early B star, which has been detected in ro-vibrational CO lines and one of
the few HAeBe stars detected in overtone CO emission. The circumstellar disk
has not been detected at mm-wavelengths, because a cold foreground cloud
absorbs all the low J 12 CO lines. These lines are typically used for
detecting gas in disk with facilities like SMA and ALMA. The MWC 297 disk does
not have any cold dust. It that sense it resembles MWC 349A, which has been
detected in CO overtone emission but has not been detected in any molecular
line at mm-wavleengths, suggesting that it has no cold molecular gas. For the
same reason we find it very unlikely that MWC 297 would have any cold
molecular gas. Based on the rovibrational lines it also looks like the disk is
rather compact. It would be very interesting to get a better estimate of the
size of the disk using spectroastrometric imaging using the rovibrational CO
lines.

We appreciate Prof. Andrea Banzatti sharing some of his iSHELL data on MWC 297
with us. We also thank Dr. Adwin Boogert for helping us with the telluric
correction of the September 2020 iSHELL spectrum. We are especially grateful
to the anonymous referee, whose constructive criticism encouraged us to
completely rethink our modeling approach to the CO lines and to get additional
data.

## References

  * Acke & van den Ancker (2004) Acke, B., & van den Ancker, M. E., et al. 2004, A&A, 426, 151
  * Acke et al. (2008) Acke, B., Verhoelst, T., van den Ancker, M. E., et al. 2008,A&A, 485, 209
  * Alonso-Albi et al. (2009) Alonso-Albi, T., Fuente, A., Bachiller, R., et al. 2009, A&A, 497, 117
  * Banzatti et al. (2022) Banzatti, A., Abernathy, K. M., Brittain, S., et al. 2022, ApJ, 163, 174
  * Banzatti et al. (2018) Banzatti, A., Garufi, A., Kama, M., et al. 2018, A&A, 609, L2
  * Banzatti & Pontoppidan (2015) Banzatti, A., Pontoppidan, K. M. 2015, ApJ, 809, 167
  * Bast et al. (2011) Bast, J. E., Brown, J.. M., Herczeg, G. J., et al. 2011, A&A, 527, A119
  * Beltrán & de Wit (2016) Beltrán, M. T., & de Wit, W. J. 2016, A&A Rev., 24, 6
  * Bik, Kaper & Waters (2006) Bik, A., Kaper, L., & Waters, L. B. F. M. 2006, A&A, 455, 561
  * Bik & Thi (2004) Bik, A., & Thi, W. F. 2004, A&A, 427, L13
  * Blake &Boogert (2004) Blake, G. A., & Boogert, A. C. A. 2004, ApJ, 606, L73
  * Blum et al. (2004) Blum, R. D., Barbosa, C. L., Damineli, A., et al. 2004, ApJ, 617, 1167
  * Bohlin et al. (1978) Bohlin, R. C., Savage, B. D., & Drake, J. F. 1978, ApJ, 224, 132
  * Boogert (2022) Boogert, A. 2022, private comm.
  * Bosman et al. (2019) Bosman, A. D., Banzatti, A., Bruderer, S., et al. 2019, A&A, 631, A133
  * Brittain, Carr & Najita (2018) Brittain, S. D., Carr, J. S., & Najita, J. R. 2018, PASP, 130, 074505
  * Brittain, Najita & Carr (2009) Brittain, S. D., Najita, J. R., & Carr, J. S. 2009, ApJ, 70, 85
  * Brittain et al. (2007) Brittain, S. D., Simon, T., Najita, J.R., et al. 2007, ApJ, 659, 685
  * Brittain et al. (2003) Brittain, S. D., Rettig, T. W., Simon, T., et al. 2003, ApJ, 588, 535
  * Brown et al. (2013) Brown, J. M., Pontoppidan, K. M., van Dishoeck, E. F., et al. 2013, ApJ, 770, 94
  * Cantó et al. (1984) Cantó, J., Rodríguez, L. F., Calvet, N., et al. 1984, ApJ, 282, 631
  * Chandra, Maheshwari & Sharma (1996) Chandra, S., Maheshwari, V. U., & Sharma, A. K. 1996, A&AS, 117, 557
  * Cooper et al. ( 2013) Cooper, H. D. B., Lumsden, S. L., Oudmaijer, A. D., et al. 2013, MNRAS, 430, 1125
  * Cowley et al. (2012) Cowley, C. R., Hubrig, S., Castelli, F., et al. 2012, A&A, 537, L6
  * Cushing, Vacca & Rayner (2004) Cushing, M. C., Vacca, W. D., & Rayner, J. T., 2004, PASP, 116, 362
  * Doppmann et al. (2017) Doppmann , G. W., Najita, J. R., Carr, J. S. 2017, ApJ, 836, 242
  * Drew et al. (1997) Drew, J. E., Busfield, G., Hoare, M. G., et al. 1997, MNRAS, 286, 538
  * Espaillat et al. (2008) Espaillat, C., Muzerolle, J., Hernández, J., et al. 2008, ApJ, 689, L145
  * Fairlamb et al. (2015) Fairlamb, J. A., Oudmaijer, R. D., Mendigutía, I., et al. 2015, MNRAS, 453, 976
  * Goto et al. (2003) Goto, M., Usuda, T., Takato, N., et al. 2003, aap, 598, 1038
  * Guzmán-Días et a. (2021) Guzmán-Días, J., Medigutía, I., Montesinos, B., et al. 2021, A&A, 650, A182
  * Haisch, Lada & Lada (2001) Haisch, K. E., Lada, E. A., & Lada, C. J. 2001, ApJ, 533, L153
  * Herczeg et al. (2011) Herczeg, G. J., Borwn, J. M., van Dishoeck, E. F., et al. 2011, A&A, 533, A112
  * Hone et al. (2017) Hone, E., Kreplin, A., Hofmann, K.-H., et al, 2017, A&A, 609, A17
  * Ilee et al. (2018) Ilee, J. D., Oudmaijer, R. D., Wheelright, H. E., et al. 2018, MNRAS, 477, 3360
  * Ilee et al. (2014) Ilee, J. D., Fairlamb, J., Oudmaijer, R. D., et al. 2014, MNRAS, 445, 3723
  * Ilee et al. (2013) Ilee, J. D., Wheelright, H. E., Oudmaijer, R. D., et al. 2013, MNRAS, 429, 2960
  * Ishii et al. (2001) Ishii, M., Nagata, T., Sato, S., et al. 2001, AJ, 121, 3191
  * Jensen et al. (2021) Jensen Jr., S. K., Brittain, S. D., Najita, J. R., et al. 2021, PASP, 133, 104402
  * Lazareff et al. (2017) Lazareff, B., Berger, J.-P., Kluska, J., et al. 2017, A&A, 599, A85
  * Kluska et al. (2020) Kluska, J., Berger, J.-P, Malbet, F., et al. 2020, A&A, 636, A116
  * Kraus et al. (2020) Kraus, M., Arias, M. L., CZidale, L.S., et al. 2020, MNRAS, 493, 4308
  * Kraus et al. (2008) Kraus, M., Hofmann, K.-H., Benisty, M., et al. 2008, A&A, 489, 1173
  * Kraus et al. (2000) Kraus, M., Krügel, E., Thum, C., et al. 2000, A&A, 362, 158
  * Malbet et al. (2007) Malbet, F., Benisty, M., de Wit, W.-J., et al. 2007, A&A, 464, 43
  * Manoj et al. (2007) Manoj, P., Ho, P. T. P., Ohashi, N., et al. 2007, ApJ, 667, L187
  * Meeus et al. (2001) Meeus, G., Waters, L. B. F. M., Bouwman, J., et al. 2001, A&A, 365,476
  * Moscadelli et al. (2021) Moscadelli, L., Cesaroni, R., Beltrán, M. T., et al. 2021, A&A, 650, A142
  * Najita et al. (2007) Najita, J. R., Carr, J. S., Glassgold, A. E., et al. 2007, in Protostars and Planets V, ed. B. Reipurth, D. Jewitt, & K. Keil (Tucson: Univ. of Arizona), 507
  * Najita, Carr & Mathieu (2003) Najita, J., Carr, J. S., & Mathieu, R. D. 2003, ApJ, 589, 931
  * Navarete et al (2021) Navarete, F., Damineli, A., Steiner, J.E., et al. 2021, MNRAS, 503, 270
  * Pomohaci et al. ( 2017) Pomohaci, R., Oudmaijer, A. D., Lumsden, S. L., et al. 2017, MNRAS, 472, 3624
  * Pontoppidan et al. (2011) Pontoppidan, K. L., Blake, G. A., & Smette, A. 2011, ApJ, 733, 84
  * Pontoppidan et al. (2008) Pontoppidan, K. L., Blake, G. A., van Dishoeck, E. F., et al. 2008, ApJ, 684, 1323
  * Rayner et al. (2022) Rayner, J., Tokunaga, A., Jaffe, D., et al. 2022, PASP, 134, 015002
  * Riello et al. (2021) Riello, M., De Angelim F., Evans, D. W., et al. 2021, A&A, 649, A3
  * Sallum et al. (2022) Sallum, S., Eisner, J. A., Stone, J. M., et al. 2022, ApJ, 161, 28
  * Salyk et al. (2009) Salyk, C., Blake, G. A., Boogert, A. C. A., et al. 2009, ApJ, 699, 330
  * Sandell et al. (2023) Sandell, G., Güsten, R., Friberg, P., et al. 2023, in preparation
  * Sandell et al. (2020) Sandell, G., Wright, M., Güsten, R., et al. 2020, ApJ, 904, 139
  * Sandell, Weintraub & Hamidouche (2011) Sandell, G., Weintraub, D. A., & Hamidouche, M. 2011, ApJ, 727, 26
  * Scoville et al. (1983) Scoville, N. Z., Kleinmann, S. G., Hall, D. N.B., et al. 1983, ApJ, 275, 201
  * Scoville et al. (1979) Scoville, N. Z., Hall, D. N.B., Kleinmann, S. G., et al. 1979, ApJ, 232, L121
  * Tuthill et al. (2002) Tuthill, P. G., Monnier, J. D., Danchi, W. C., et al. 2002, ApJ, 577, 826
  * Vacca &Sandell (2022) Vacca, W. D., & Sandell, G. 2022, ApJ, 941, 189
  * Vacca, Cushing & Rayner (2003) Vacca, W. D., Cushing, M. C. & Rayner, J. T. 2003, PASP,115, 389
  * van der Plas et al. (2015) van der Plas, G., van der Ancker, M. E., Waters, L. B. F. M., et al. 2015, A&A, 574, A75
  * Vioque et al. (2018) Vioque, M., Oudmaijer, R. D., Baines, D., et al. 2018, A&A, 620, A128
  * Weigelt et al. (2011) Weigelt, G., Grinin, V. P., Groh, J. H., et al. 2011, A&A, 527, A103
  * Wheelwright et al. (2010) Wheelwright, H. E., Oudmaijer, R. D., de Wit, W. J., et al. 2010, MNRAS, 408, 1840
  * Williams & Cieza (2011) Williams, J. P., & Cieza, L. A. 2011, ARA&A, 49, 67
  * Wilson & Rood (1994) Wilson, T. L., & Rood, R. T. 1994, ARA&A, 32, 191
  * Zapata et al. (2019) Zapata, L. A., Garay, G., Palau, A., et al. 2019, ApJ, 872, 176