A Rotating Accretion Disk around MWC 297, a Young B1.5Ve Star

High-resolution spectra with iSHELL on IRTF in the K and M bands of the young, heavily accreting B1.5Ve 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 shows 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 2–1. We see strong absorption in 12CO and 13CO 1–0 at low J levels, which is due to 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.


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 ∼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 gsandell@hawaii.edubill.vacca@noirlab.edustellar 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 µm and at 4.5 -.2 µm.CO overtone emission (υ = 2 -0 and υ = 3 -1) was first detected in the BN object (mass ∼ 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 ∼ 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 (υ = 1 -0) rovibrational lines, located in the M band, 4.5 -5.2 µm, 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 ∼ 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 > µm.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 υ = 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 µm as well as the rovibrational spectra in the M-band.
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 crossdispersed 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 µm 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. ′′ 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. ′′ 375 slit spans ∼ 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. ′′ 375, which yields a nominal resolving power of 88,000 for the spectra.At the distance of MWC 297, 418 pc, the iShell 0. ′′ 375 slit spans ∼ 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 µm.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 to-gether 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 modelcorrected 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 ≲ 15, while it was +41.2 km s −1 in 2022, which largely shifted all CO lines away from telluric absorption features.

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 (∆υ = 1) transitions of CO at 4.6 µm.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 (≳ 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 (≲ 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 µm 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 µm 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. 2018Banzatti et al. , 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, ∼ 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 ≳ 2.5 are triangular, and ≲ 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 υ = 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. (2018Banzatti et al. ( , 2022)); Bosman et al. (2019) suggests that one can identify three emitting regions that show a dependence on the υ2/υ1 ratio depending on the location of their emitting radii, R CO , compared to their dust sublimation radius, R subl .Regions a and 1b both have high vibrational ratio υ2/υ1 = 0.2 -.0, but come from very different regions in the disk, although both are largely dust free.Region 1a is inside R 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 υ2/υ1 = 0.01 -0.2, and the CO is emitted from regions outside the dust rim.

ANALYSIS AND RESULTS
The iSHELL spectrum exhibits a large number of double-peaked emission lines from rovibrational transitions of CO.The υ = 1 -0 CO lines up to rotational level P(47) and higher vibrational levels up to υ = 5 -4 are clearly seen.Emission from 13 CO is also seen, in both the υ = -0, and υ = 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 β, Hu ϵ, and Hu δ) 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.
At 10% level the line widths are similar or ∼ 65 km s −1 , suggesting that the inner radius is similar for all of them.The FW10% maybe slightly larger for the 12 CO 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 12 CO 1 -0 than what we find, 45.3 km s −1 .We determine a V lsr = +6.7 ± 2.2 km s −1 for 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.

Modeling the emission lines with a rotating
Keplerian disk.Inspection of the K and M band spectra yields conflicting information regarding the CO emitting gas.The relative strengths of the 12 CO lines across the M band suggest that the gas is optically thick.In addition, the 13 CO emission lines are much stronger relative to those of 12 CO than expected given the normal isotopic ratio and optically thin gas.However, the clearly doublepeaked 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 ⊙ 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 12 CO/ 13 CO 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 K and a density of 1.6 × 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 13 CO abundance needs to be increased by a factor of 6 -7 compared to normal interstellar abundances.
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 ∼ 88, 000 in the M band), and sampled it at the observed wavelength sampling of the observations.A χ 2 comparison between the observed 12 CO 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 13 CO 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 13 CO emission line strengths from the model and that of the data on the rotational diagram, although the slope of the points for the 13 CO 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 12 CO and 13 CO emission profiles, suggests that the 12 CO and 13 CO emission arises from different parts of the disk.Therefore, we generated a pure 13 CO emission model in which the 13 CO gas temperature was the same as that for the 12 CO gas, but the abundance was enhanced by a factor of 6 , which gave a good fit to the 13 CO 1 -0 spectra.Figure 5 shows the observed spectra overplotted with the best fit model without any additional scaling for 13 CO.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 12 CO 2-1 lines in the model are much stronger than in the data.In addition, the line profiles of many of the low lying 12 CO -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.

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 13 CO 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 ∼ 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 υ = 1 -0 12 CO emission 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 12 CO and at the bottom shows 13 CO 1 -0.Regions of strong telluric absorption have been blanked out.
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 13 CO absorption lines, while there is a clear shift in velocity at the lowest levels.We illustrate this in Figure 7, which shows three CO υ 1 -0 lines and three of the 13 CO 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 12 CO and 13 CO lines we find a velocity of -0.6 ± 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 ± 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 ∼ 7.2 km s −1 of the molecular cloud from 13 CO(6-5) observations with APEX.This agrees well with 13 CO(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 12 CO 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 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, 13 CO lines.Using all 13 CO lines in the analysis resulted in large densities (2 10 18 cm −2 ) and very low temperatures (∼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  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. lines.As can be seen in Figure 7, even the 13 CO 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 ∼ 6.5 10 17 cm −2 and a temperature of ∼ 8.5 K.We then combined the two weakest CO lines with the two weakest 13 CO 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 −1 for this fit, which is in aqreement with our resolution (R ∼ 88,000).This resulted in a CO column density, N(CO) = 6.7 ± 0.4 10 17 cm −2 , and a temperature T = 8.25 ± 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 13 CO.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.

DISCUSSION
Generally half of all solar mass stars have lost their disks after 3 Myr with an overall disk life time of 6 Myr Figure 7. 12 CO and 13 CO υ = 1 -0 absorption lines.The 12 CO profiles are plotted with solid lines and 13 CO 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 12 CO υ = 1 -0 R(0) line appears flat topped, much broader, and shifted in velocity relative to 12 CO 1 -0 R(4) and 13 CO 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 12 CO 1 -0 R(0) line.It is also seen in the 13 CO 1 -0 line as a strong red-shifted wing.The 0 km s −1 velocity corresponds to a V lsr = 6.7 km s −1 .(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 LkHα 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 ∼ 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 LkHα 101, as it is still heavily accreting with an age in the range of -2.8 × 10 5 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 re- iSHELL spectrum showing the υ = 1 -0 12 CO 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 β and Hu ϵ are also seen.The model fit to the absorption lines is overlaid in red.
ally 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. (2018Banzatti et al. ( , 2022)); Bosman et al. (2019).It has double peaked profiles and the ratio of vibrational levels of υ2/υ1 = 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 υ = 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(Ilee et al. , 2014) ) with gas temperatures of ∼ 3000 K and column densities of 10 21 cm −2 .The overtone emission in MWC 297 is much colder, ∼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 υ = 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 υ = 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 υ = 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 (Brγ) 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 ∼ 5 AU1 .As pointed out by Vacca &Sandell (2022) the Brγ 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, ∼ 65.2 km s −1 , we find the inner CO radius ∼ 5.6 AU, i.e. roughly at the same radius as the dust destruction radius, A normal isotope ratio of [ 12 C/ 13 C] of 77 greatly underestimates the strength of the observed 13 CO line intensities.We find that a relative 12 CO to 13 CO abundance of ∼ 6 provides a good fit to the data.This does not mean that the disk has anomalous isotope ratios, but rather that 13 CO is not in LTE and coexisting with 12 CO.As we seen, the FWHM of 13 CO 1 -0 differs by more than 5% from that of 1 -0, see Section 4. These kind of 12 CO/ 13 CO ratios are not uncommon.van der Plas et al. (2015) found that the 12 CO/ 13 CO ratio varied between and 7 for their sample of 12 HAeBe stars and that the rotational temperatures of 13 CO were on the average lower than those for 12 CO.Banzatti et al. (2022), who has the largest sample of rovibrational CO lines in HAeBe stars, do not quote a 12 CO/ 13 CO 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 [ 12 CO/ 13 CO], because for most disks 12 CO 1 -0 is more optically thick than 13 CO 1 -0 , nor does it account for difference in excitation between 12 CO and 13 CO.
We detect narrow absorption lines, FWHM ∼ 4 -5 km s −1 from a cold unrelated foreground cloud in both 12 CO and 13 CO.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 12 CO are completely saturated.At higher J levels the lines become more blue-shifted and agree in velocity with the 13 CO absorption lines.The lowest transitions of 13 CO is also optically thick and show a red-shifted line wing.wing.Detailed analysis shows that the redshifted component comes from a colder, but less dense cloud component.For the blue-shifted cloud component, which is at a V lsr ∼ 6.4 km s −1 we derive a temperature of 8.3 ± 0.1 K and a CO column density, N(CO) = 6.7 × 10 17 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 .3magnitude 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 υ = 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 13 CO.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 12 CO at the lowest J levels.

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 12 CO υ = 1-0 emission is optically thick for the low excitation lines.Even the 13 CO 1 -0 and 12 CO 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 × 10 18 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 12 CO 2-1 lines in the model are much stronger than in the data.We also find that a normal isotope ratio of [ 12 C/ 13 C] of 77 greatly underestimates the 13 CO line intensities, while a 12 CO to 13 CO abundance of ∼ 13 provides a good fit to the data.Other HAeBe disks typically show similarly enhanced 13 CO line intensities, most likely because excitation conditions in the disks differ between 12 CO and 13 CO.
The 12 CO 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 13 CO 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 × 10 17 cm −2 , which correspond to a visual extinction of 6.3 mag.With our velocity resolution we could not separate the second redshifted 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 mmwavleengths, 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.

Figure 2 .
Figure2.Median line profiles of 12 CO υ = 1 -0 and 2 -1 and 13 CO -0 and 12 CO 2 -0 emission lines.The dashed blue line is the 12 CO 2 -1 profile scaled by 2.25, the dashed red line is the 13 CO 1 -0 profile scaled by 3, the dashed orange is the scaled 12 CO 2 -0.The scaled profiles demonstrate the differences in widths compared to the 12 CO -0 profile.

Figure 3 .
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 .
Figure 4.The ring model for a rotational temperature of 1500

Figure 6 .
Figure 6.An example illustrating the extraction of absorption line profiles.The observed emission line, in this case the 12 CO υ = 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.

Table 1 .
Line profile parameters for median averaged emission lines.The FWHM values have measurement errors of ∼ 0.5 km s −1 , the errors for the peak separations are ∼ 1 km s −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 µm is very weak.