Investigating the Impact of Metallicity on Star Formation in the Outer Galaxy.

I. VLT/KMOS Survey of Young Stellar Objects in Canis Major
Agata Karska

0000-0002-9051-1781]Dominika Itrich
Institute of Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziądzka 5, 87-100 Toruń, Poland
European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching bei München, Germany
Universitäts-Sternwarte, Ludwig-Maximilians-Universität, Scheinerstrasse 1, 81679 München, Germany

0000-0001-8913-925X]Agata Karska
Institute of Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziądzka 5, 87-100 Toruń, Poland
Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn, Germany

0000-0003-2248-6032]Marta Sewiło
Exoplanets and Stellar Astrophysics Laboratory, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA
Center for Research and Exploration in Space Science and Technology, NASA Goddard Space Flight Center, Greenbelt, MD 20771
Department of Astronomy, University of Maryland, College Park, MD 20742, USA

Lars E. Kristensen
Niels Bohr Institute, Centre for Star and Planet Formation, University of Copenhagen, Øster Voldgade 5-7, 1350 Copenhagen K, Denmark

Gregory J. Herczeg
Kavli Institute for Astronomy and Astrophysics, Peking University, Yiheyuan 5, Haidian Qu, 100871 Beijing, China
Department of Astronomy, Peking University, Yiheyuan 5, Haidian Qu, 100871 Beijing, China

Suzanne Ramsay
European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching bei München, Germany

0000-0002-3747-2496]William J. Fischer
Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA

Benoît Tabone
Leiden Observatory,Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands
Institut d'Astrophysique Spatiale, Université Paris-Saclay, CNRS,  B$\hat{a}$timent 121, 91405 Orsay Cedex, France

Will R. M. Rocha
Laboratory for Astrophysics, Leiden Observatory,Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands

0000-0001-5785-1154]Maciej Koprowski
Institute of Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziądzka 5, 87-100 Toruń, Poland

Ngân Lê
Institute of Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziądzka 5, 87-100 Toruń, Poland

Beata Deka-Szymankiewicz
Institute of Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziądzka 5, 87-100 Toruń, Poland

The effects of metallicity on the evolution of protoplanetary disks may be studied in the outer Galaxy where the metallicity is lower than in the solar neighborhood. We present the VLT/KMOS integral field spectroscopy in the near-infrared of $\sim$120 candidate young stellar objects (YSOs) in the CMa-$\ell$224 star-forming region located at a Galactocentric distance of 9.1 kpc.
We characterise the YSO accretion luminosities and accretion
rates using the hydrogen Br$\gamma$
emission and find the median accretion luminosity of $\log{(L_{\rm acc})} = -0.82^{+0.80}_{-0.82} L_\odot$.
Based on the measured accretion luminosities, we investigate the hypothesis of star formation history in the CMa-$\ell$224. Their median values suggest that Cluster C, where most of YSO candidates have been identified, might be the most evolved part of the region.
The accretion luminosities are similar to those observed toward low-mass YSOs in the Perseus and Orion molecular clouds, and do not reveal the impact of lower metallicity.
Similar studies in other outer Galaxy clouds covering a wide range of metallicities are critical to gain a complete picture of star formation in the Galaxy.

§ INTRODUCTION

Stars form as a result of complex physico-chemical processes initiated by the gravitational collapse of a dense and cold molecular cloud. The formation of a rotating envelope and the embedded disk is associated with the ejection of jets and disk winds, together responsible for the removal of the angular momentum <cit.>. The interaction between jets/winds and the envelope leads to the formation of outflow cavities and the dispersion of some mass reservoir [van Kempen et al., 2010, Visser et al., 2012].  Shock waves at the outflow/envelope interface compress and heat the envelope material to hundreds or thousands of K, even around low-mass protostars [Kristensen et al., 2017, Karska et al., 2018].

The net mass growth of a young star is a balance between the mass accretion from the envelope-disk system and the mass ejection by jets and winds. The main accretion phase occurs during the earliest evolutionary stages of a young stellar object (YSO; Class 0 and Class I), accompanied by collimated H$_2$ jets [Davis & Eisloeffel, 1995, Stanke et al., 2002, Kristensen et al., 2007] and more extended molecular outflows [van der Marel et al., 2013, Tobin et al., 2016, Mottram et al., 2017]. Once the protostar evolves into Class II, the envelope mass reservoir is mostly depleted, the accretion rate decreases, and the outflow opening angle widens [Offner et al., 2011, Agra-Amboage et al., 2014]. The gradual decrease of accretion rates  [Manara et al., 2012, Ansdell et al., 2017, Testi et al., 2022] is accompanied by decreasing mass loss rates and a transition from mostly molecular to atomic/ionic outflows [Nisini et al., 2015, Bally, 2016].

Herschel/SPIRE 250 $\mu$m image of the main filament in the CMa-$\ell$224 star-forming region. Grey circles show the positions of YSO candidates observed with VLT/KMOS, while light-grey dashed ellipses indicate regions hosting clusters of sources with the extended 4.5 $\mu$m emission (most likely tracing outflows) defined in <cit.>. North is up, east to the left.

The location of a star-forming region in the Galaxy may influence the process of mass assembly. In the outer Galaxy, the gas surface density in molecular clouds is known to be lower than those in the solar neighborhood [Roman-Duval et al., 2010]. The decrease of metallicity with Galactocentric radius, traced by Fe or O abundance gradients, also translates to lower dust and molecular gas abundances  [Sodroski et al., 1997, Lépine et al., 2011, Hawkins, 2022]. Additionally, lower cosmic-ray fluxes and the interstellar UV radiation field reduce the amount of gas and dust heating [Bloemen et al., 1984]. These factors may affect physical and chemical conditions in star forming regions and result in globally lower star formation rates and efficiencies in the outer Galaxy [Kennicutt & Evans, 2012, Heyer & Dame, 2015, Djordjevic et al., 2019].

Recent infrared studies revealed significant star formation activity in the Canis Major star-forming region in the outer Galaxy [Fischer et al., 2016, Sewiło et al., 2019].
The 2.$^{\circ}5\times1.^{\circ}4$ region in Canis Major dubbed CMa-$\ell224$ centered on ($l$, $b$)=(224.$^\circ$5, $-$0.$^\circ$65) is particularly interesting because it hosts clusters of sources with extended 4.5 $\mu$m emission (likely tracing outflows), many of which were identified as YSOs [Sewiło et al., 2019]. CMa-$\ell224$ is located at a distance of 0.92 kpc from the Sun and 9.1 kpc from the Galactic Center (e.g., Claria, 1974),
where a subsolar metallicity is predicted by the O/H Galactocentric radial gradients determined based on the observations of H ii regions ($Z$$\sim$$0.55-0.73$ $Z_\odot$: Balser et al., 2011, Fernández-Martín et al., 2017, Esteban & García-Rojas, 2018). CMa-$\ell224$ corresponds to a  local peak of the H$_2$ column density and the lowest  dust temperatures in the far-infrared observations of the outer Galaxy by the Herschel infrared Galactic Plane Survey” (Hi-GAL; Molinari et al., 2010).  The region contains a  complex network of star-forming filaments <cit.> and a widespread emission in CO and  its isotopologues,  tracing dense, molecular gas [Olmi et al., 2016, Benedettini et al., 2020, Lin et al., 2021].

[Sewiło et al., 2019] identified 294 YSO candidates in CMa-$\ell$224 based on the photometric data from the near- and mid-infrared catalogs: Spitzer's GLIMPSE360: Completing the Spitzer Galactic Plane Survey” (PI: B. Whitney), the Two Micron All Sky Survey (2MASS; Skrutskie et al., 2006), and the AllWISE catalog that combines the Wide-field Infrared Survey Explorer” (WISE; Wright et al., 2010) and NEOWISE [Mainzer et al., 2011] data. Some YSO candidates in CMa-$\ell$224 are associated with the extended 4.5 $\mu$m emission (Extended Green Objects, EGOs; Cyganowski et al., 2008), likely dominated by the H$_2$ line emission from outflow shocks (e.g., Cyganowski et al., 2011). The Spectral Energy Distribution (SED) fitting with the [Robitaille, 2017] YSO models identified 37 sources with envelopes, i.e., Class 0/I YSOs. Thus,  CMa-$\ell$224 offers an opportunity to evaluate how metallicity affects the ongoing star formation in a region that is significantly closer than, for example, the Magellanic Clouds.

In this paper, we address the following questions.
What are the mass accretion
rates in YSOs in CMa-$\ell$224? How are they connected to the star formation scenarios in this region? Is there any impact of the reduced gas metallicity on the accretion
properties of YSOs in CMa-$\ell$224?

To this end, we present the results of the  integral field spectroscopy of 124 YSO candidates  in the CMa-$\ell$224 star-forming region using the $K$-band Multi Object Spectrograph (KMOS; Sharples et al., 2013) on the Very Large Telescope (VLT). VLT/KMOS provides $2\rlap.{''}8\times2\rlap.{''}8$ spectral maps in ro-vibrational H$_2$ lines, hydrogen Br$\gamma$, and CO bandhead at 2.3 $\mu$m. At the distance of CMa-$\ell$224, we obtained $\sim$$2580\times2580$ au maps at the physical resolution of 184 au.

The paper is organized as follows. In Section 2, we describe our sample selection, KMOS observations, and data reduction. In Section 3, we present the continuum and line emission maps, provide the statistics on line detections, discuss gas spatial distribution and describe calculations of the
mass accretion rates. In Section 4, we discuss the results in the context of the star formation scenarios in CMa-$\ell$224, and in Section 5, we present the summary and conclusions.

This paper is the first in a series presenting the multi-wavelength spectroscopy of YSO candidates in the CMa-$\ell$224 star-forming region. Two forthcoming papers will discuss the $^{13}$CO and C$^{18}$O 2-1 observations of YSO candidates in the main filament in CMa-$\ell$224 with the Atacama Large Millimeter/submillimeter Array (ALMA; Koprowski et al., in prep.), and the spectral types and excess continuum measurements toward YSO candidates in the second brightest filament in the region using the SpeX instrument at the NASA Infrared Telescope Facility (IRTF; Le et al., in prep.).

§ KMOS OBSERVATIONS AND DATA REDUCTION

We used VLT/KMOS to observe YSO candidates in the CMa-$\ell$224 star-forming region as part of the ESO programme 0102.C-0914(A) (PI: A. Karska). The sources were selected from the catalogue of YSO candidates from [Sewiło et al., 2019] based  on their 2MASS $K$-band brightness and location within two KMOS patrol fields.

Targets were selected from the $K$-band magnitude range of 9.5–15 to ensure a sufficient signal-to-noise ratio (SNR) without saturation for a single integration time applied to all exposures.

KMOS is a second-generation instrument operating in the near-infrared (near-IR) at the B Nasmyth focus of the VLT Unit Telescope 1 on Paranal in Chile, in operation since November 2012 [Sharples et al., 2013]. It performs Integral Field Spectroscopy (IFS) for up to 24 targets simultaneously. Each of the 24 pick-off arms is connected to one of three identical spectrographs and detectors. KMOS arms are allocated in two planes (top” and bottom”) to avoid interference between them. The patrol field is 7$\rlap.{'}$2 in diameter, the size of each IFU is $2\rlap.{''}8\times2\rlap.{''}8$, and the spatial sampling (spaxel size) is $0\rlap.{''}2\times0\rlap.{''}2$.

Observations were prepared using the KMOS Arm Allocator (KARMA[<https://www.eso.org/sci/observing/phase2/SMGuidelines/KARMA.html>]) and p2[<https://www.eso.org/sci/observing/phase2/p2intro.html>], a web-based tool for preparation of the Phase 2 material. The observations were performed in October and December 2018.

We used the K grating with the spectral range from 1.93 to 2.50 $\mu$m and  resolving power $R$ of $\sim$4000, with a spectral sampling $\Delta\lambda$ of $\sim$2.8 Å per spectral element, corresponding to $\sim$39.7 km s$^{-1}$ at 2.1218 $\mu$m. The total integration time for each pointing was 1300 s with 5 single exposures of 260 s, dithered by 0$\rlap.{''}$2.

The Nod-to-Sky mode was used to observe sources in both the science and sky pointings.  The arms originally observing targets were next directed to the sky area, while those observing the sky were capturing science objects in the sky pointing, usually offset by a few arcmin. One of the pick-off arms was broken during semester 102, which resulted in a lack of observations of 6 sources.

Table <ref> shows the summary of the atmospheric conditions during our KMOS observations. The amount of precipitatable water vapour (PWV) was in the range from $\sim$2.5 to 4 mm, and the sky was covered with thin cirrus (TN). Seeing measured at the observatory site was typically below 1$''$ with the point spread function (PSF) of the data in the 0$\rlap.{''}$4–0$\rlap.{''}$7 range.

Atmospheric Conditions

OB      Date        Seeing     2cAirmassa     PWV     Sky     Grade

(”)       SCI     STD     (mm)

1     23.10.2018     $\sim$1.5     1.20     1.70     $<$2.5        TN     B

2     24.10.2018     $\sim$0.5     1.18     1.15     $\sim$2.5     TN     B

3     25.10.2018     $\sim$0.6     1.25     1.83     $\sim$2.5     TN     B

4     11.12.2018     $\sim$1.0     1.34     1.20     $\sim$4.0     TN     B

5     11.12.2018     $<$0.6        1.13     1.20     $\sim$4.0     TN     B

6     11.12.2018     $\sim$0.5     1.05     1.20     $>$3.0        TN     B

aA mean airmass during the integration made for one of the targets per OB (SCI) and for a standard star used for telluric correction (STD).

Data reduction was performed using the ESO Recipe Execution Tool (, version 2.0.2 of the KMOS pipeline), the terminal-based software. The standard procedure included the processing of the raw calibration data: dark, flat-field, illumination,  and sky subtraction (a single sky observation per science target). The wavelength calibration was done using the Argon – Neon lamp.

We used a standard procedure for the telluric correction and flux calibration by observing a standard star, one per each of the three detectors, and comparing observations to the stellar model.
The standard pipeline does not account for the difference in the airmass between the science and standard star observations, which might lead to under- or over-estimation of the telluric lines of up to 10%. Additionally, each of the 24 arms have a slightly different spectral resolution resulting in different line shapes.
It is possible to account for these effects by modeling the telluric lines using the science or the standard star spectrum [Coccato et al., 2019]. However, for our science data with the median SNR measured on the featureless parts of the spectra of less than 50, the differences between the standard method we used and modelling are negligible since the noise dominates any uncertainty coming from the telluric standard [Coccato et al., 2019].

The imperfect telluric correction is a source of the unknown uncertainty, especially in the red part of the $K$-band (above 2.4 $\mu$m), where some of the H$_2$ lines are located. Also, other telluric features are poorly corrected, e.g., near 2 $\mu$m. All spectral regions particularly affected by the telluric lines are marked on the figures.

The reduced  single exposures were collapsed into the final data cubes. The entire process of creating the 3D KMOS science data cubes is described in [Davies et al., 2013]. For the analysis of sources observed multiple times, we used the data cubes combined with the KMOS pipeline task combine.

The $K$-band continuum fluxes are calculated by fitting a 7th order Chebyshev polynomial to the spectrum in each spaxel (spatial pixel) of the data cube. The continuum level is defined as the value of the fit at 2.12 $\mu$m, near the middle of the Johnson $K$-band and the rest wavelength of the 1-0 S(1) H$_2$ line. The noise of the spectrum is calculated as the root mean square (RMS) of deviations from the continuum in the line-free parts of the spectrum: 2.076–2.106 $\mu$m, 2.125–2.143 $\mu$m, 2.179–2.187 $\mu$m, 2.198–2.201 $\mu$m, and 2.268–2.274 $\mu$m. In four cases of faint objects (No. 11, 14, 26, 69), we integrate the continuum emission over the full spectral range in each spaxel to detect these sources at a $>$3$\sigma$ level and verify their coordinates; the spectral analysis remained unchanged.

Catalog of Sources in CMa-$\ell$224 Detected with VLT/KMOS

IRAC Designationa
Continuum near 2.12 $\mu$m

($^h$ $^m$ $^s$)
($^o$ $\arcmin$ $\arcsec$)
(10$^{-17}$ erg s$^{-1}$ cm$^{-2}$ Å$^{-1}$)

1      	G224.2025-00.8569	   	 07 09 07.09 	   	 -10 24 47.50 	   	 08.27 	$\pm$	 0.12    	 II 	   	 d

2  	   	G224.2265-00.8620	   	 07 09 08.69 	   	 -10 26 12.22 	   	 26.40 	$\pm$	 0.13    	 II 	   	 e+d 	   	 ext. H$_2$ emission

3  	   	G224.2420-00.8313	   	 07 09 17.08 	   	 -10 26 10.67 	   	 13.05 	$\pm$	 0.14    	 II 	   	 d

4  	   	G224.2449-00.9307	   	 07 08 55.89 	   	 -10 29 04.97 	   	 07.02 	$\pm$	 0.21    	 I 	   	 null 	   	 ext. H$_2$ emission

5  	   	G224.2470-00.9334	   	 07 08 55.53 	   	 -10 29 16.20 	   	 38.74 	$\pm$	 0.60    	 I 	   	 e

6  	   	G224.2483-00.9176	   	 07 08 59.12 	   	 -10 28 54.25 	   	 11.22 	$\pm$	 0.14    	 II 	   	 d

7  	   	G224.2512-00.9297	   	 07 08 56.82 	   	 -10 29 23.57 	   	 07.49 	$\pm$	 0.29    	 I 	   	 null

8  	   	G224.2530-00.8250	   	 07 09 19.67 	   	 -10 26 35.06 	   	 07.84 	$\pm$	 0.15    	 II 	   	 d

9  	   	G224.2535-00.8306	   	 07 09 18.56 	   	 -10 26 45.89 	   	 10.17 	$\pm$	 0.10    	 II 	   	 d

10 	   	G224.2567-00.8346	   	 07 09 18.06 	   	 -10 27 02.03 	   	 11.21 	$\pm$	 0.13    	 II 	   	 d 	   	 near edged

11c 	   	G224.2591-00.8305	   	 07 09 19.19 	   	 -10 27 03.71 	   	 2.52 	$\pm$	 0.57	   	 I 	   	 null 	   	 integrated cont.

12  	   	G224.2598-00.8398	   	 07 09 17.27 	   	 -10 27 20.77 	   	 28.52 	$\pm$	 0.13	   	 II 	   	 d 	   	 near edge

13 	   	G224.2621-00.8470	   	 07 09 15.97 	   	 -10 27 39.63 	   	 13.70 	$\pm$	 0.08    	 I 	   	 d 	   	 near edge

14c   	   	G224.2638-00.8456	   	 07 09 16.47 	   	 -10 27 43.32 	   	 2.88 	$\pm$	 0.38    	 I 	   	 null 	   	 integrated cont.

15  	   	G224.2652-00.8647	   	 07 09 12.46 	   	 -10 28 20.32 	   	 13.27 	$\pm$	 0.12    	 I 	   	 d 	   	 ext. H$_2$ emission

16  	   	G224.2674-00.8609	   	 07 09 13.54 	   	 -10 28 21.20 	   	 31.76 	$\pm$	 0.17    	 II 	   	 d

17  	   	G224.2680-00.8578	   	 07 09 14.27 	   	 -10 28 17.95 	   	 3.29 	$\pm$	 0.12    	 II 	   	 null

18  	   	G224.2697-00.8168	   	 07 09 23.38 	   	 -10 27 14.62 	   	 5.29 	$\pm$	 0.07    	 II 	   	 null

19  	   	G224.2724-00.8032	   	 07 09 26.64 	   	 -10 27 00.62 	   	 3.11 	$\pm$	 0.11    	 I 	   	 d 	   	near edge, ext. H$_2$ emission

20  	   	G224.2769-00.8049	   	 07 09 26.78 	   	 -10 27 17.65 	   	 20.18 	$\pm$	 0.10    	 II 	   	 d 	   	 near edge

21A 	   	G224.2778-00.8465	   	 07 09 17.76 	   	 -10 28 28.62 	   	 0.85 	$\pm$	 0.13    	 I 	   	 null 	   	 on edge

21B 	   	G224.2778-00.8465	   	 07 09 17.83 	   	 -10 28 29.02 	   	 0.44 	$\pm$	 0.12      	   	  	   	 separation 1073 au

22  	   	G224.2784-00.8408	   	 07 09 19.14 	   	 -10 28 21.54 	   	 11.88 	$\pm$	 0.06    	 II 	   	 null 	   	near edge

23  	   	G224.2800-00.8346	   	 07 09 20.64 	   	 -10 28 17.05 	   	 25.24 	$\pm$	 0.14    	 I 	   	 null

24  	   	G224.2802-00.8747	   	 07 09 12.00 	   	 -10 29 25.00 	   	 15.79 	$\pm$	 0.35    	 II 	   	 d

25  	   	G224.2802-00.9179	   	 07 09 02.64 	   	 -10 30 36.58 	   	 13.07 	$\pm$	 0.31    	 II 	   	 null

26c  	   	G224.2828-00.8148	   	 07 09 25.28 	   	 -10 27 53.42 	   	 1.07 	$\pm$	 0.08    	 I 	   	 d 	   	 integrated cont.

28  	   	G224.2872-00.8673	   	 07 09 14.41 	   	 -10 29 35.21 	   	 24.28 	$\pm$	 0.40    	 II 	   	 d

29  	   	G224.2881-00.9333	   	 07 09 00.17 	   	 -10 31 27.03 	   	 20.35 	$\pm$	 0.28    	 II 	   	 null

30  	   	G224.2882-00.8515	   	 07 09 17.92 	   	 -10 29 10.51 	   	 13.91 	$\pm$	 0.10    	 II 	   	 d 	   	 near edge

31  	   	G224.2903-00.8108	   	 07 09 27.00 	   	 -10 28 10.52 	   	 2.08 	$\pm$	 0.14    	 II 	   	 null

32  	   	G224.2905-00.8364	   	 07 09 21.45 	   	 -10 28 53.19 	   	 68.62 	$\pm$	 0.10    	 II 	   	 d

33  	   	G224.2914-00.8402	   	 07 09 20.72 	   	 -10 29 02.78 	   	 6.02 	$\pm$	 0.09    	 II 	   	 d

34  	   	G224.2928-00.8403	   	 07 09 20.84 	   	 -10 29 07.42 	   	 11.50 	$\pm$	 0.07    	 II 	   	 null

35  	   	G224.2939-00.8157	   	 07 09 26.29 	   	 -10 28 30.56 	   	 13.85 	$\pm$	 0.21    	 II 	   	 e+d

36A 	   	G224.2965-00.8191	   	 07 09 25.86 	   	 -10 28 44.29 	   	 98.54 	$\pm$	 0.41    	II	   	d	   	separation 1086 au

36B 	   	G224.2965-00.8191	   	 07 09 25.94 	   	 -10 28 44.29 	   	 6.51 	$\pm$	 0.04    	 	   	  	   	on edge

38A 	   	G224.2983-00.8203	   	 07 09 25.84 	   	 -10 28 51.25 	   	 12.54 	$\pm$	 0.10    	 I 	   	 null

38B 	   	G224.2983-00.8203	   	 07 09 25.78 	   	 -10 28 52.25 	   	 2.84 	$\pm$	 0.09      	   	 	   	 separation 1178 au

39  	   	G224.3005-00.8417	   	 07 09 21.41 	   	 -10 29 34.38 	   	 193.49 	$\pm$	 0.67    	 II 	   	 d 	   	 ext. H$_2$ emission

40  	   	G224.3006-00.8470	   	 07 09 20.27 	   	 -10 29 43.74 	   	 680.28 	$\pm$	 15.78	   	 II 	   	 d

41  	   	G224.3022-00.8221	   	 07 09 25.85 	   	 -10 29 07.35 	   	 23.01 	$\pm$	 0.10    	 I 	   	 null

42  	   	G224.3028-00.8329	   	 07 09 23.58 	   	 -10 29 27.12 	   	 14.65 	$\pm$	 0.29    	 II 	   	 d

43  	   	G224.3044-00.9038	   	 07 09 08.45 	   	 -10 31 30.57 	   	 14.57 	$\pm$	 0.17    	 II 	   	 d

44  	   	G224.3057-00.8247	   	 07 09 25.68 	   	 -10 29 23.07 	   	 131.32 	$\pm$	 0.67	   	 I 	   	 null

45  	   	G224.3063-00.8310	   	 07 09 24.37 	   	 -10 29 35.76 	   	 17.46 	$\pm$	 0.55    	 II 	   	 d

46  	   	G224.3083-00.8093	   	 07 09 29.31 	   	 -10 29 05.72 	   	 113.91 	$\pm$	 0.13	   	 II 	   	 d 	   	 ext. H$_2$ emission

47  	   	G224.3095-00.8172	   	 07 09 27.71 	   	 -10 29 22.66 	   	 3.73 	$\pm$	 0.08    	 I 	   	 e 	   	 ext. H$_2$ emission

48  	   	G224.3112-00.8536	   	 07 09 20.01 	   	 -10 30 28.89 	   	 41.21 	$\pm$	 0.05    	 II 	   	 d

49  	   	G224.3159-00.7490	   	 07 09 43.27 	   	 -10 27 50.24 	   	 13.38 	$\pm$	 0.12    	 I 	   	 d 	   	 near edge

50  	   	G224.3173-00.8757	   	 07 09 15.96 	   	 -10 31 25.09 	   	 12.63 	$\pm$	 0.07    	 II 	   	 d

51  	   	G224.3204-00.8462	   	 07 09 22.65 	   	 -10 30 45.91 	   	 29.36 	$\pm$	 0.50    	 II 	   	 null

52A 	   	G224.3215-00.8319	   	 07 09 25.88 	   	 -10 30 25.24 	   	 08.71 	$\pm$	 0.02    	 II 	   	 null

52B 	   	G224.3215-00.8319	   	 07 09 25.83 	   	 -10 30 25.24 	   	 6.85 	$\pm$	 0.02      	   	 	   	 separation 552 au

53  	   	G224.3222-00.8717	   	 07 09 17.36 	   	 -10 31 33.78 	   	 19.90 	$\pm$	 0.46    	 II 	   	 d

54  	   	G224.3225-00.8490	   	 07 09 22.28 	   	 -10 30 56.90 	   	 702.81 	$\pm$	 1.86	   	 I 	   	 null 	   	 ext. H$_2$ emission

55A 	   	G224.3233-00.7848	   	 07 09 36.31 	   	 -10 29 12.99 	   	 33.38 	$\pm$	 0.13    	 II 	   	 d

55B 	   	G224.3233-00.7848	   	 07 09 36.27 	   	 -10 29 13.19 	   	 27.61 	$\pm$	 0.08      	   	 	   	 separation 582 au

56  	   	G224.3243-00.8696	   	 07 09 18.06 	   	 -10 31 37.05 	   	 11.24 	$\pm$	 0.22    	 II 	   	 d

57  	   	G224.3268-00.8331	   	 07 09 26.22 	   	 -10 30 44.05 	   	 10.33 	$\pm$	 0.06    	 II 	   	 null

58  	   	G224.3285-00.8597	   	 07 09 20.64 	   	 -10 31 34.20 	   	 134.72 	$\pm$	 1.35	   	 II 	   	 d

59  	   	G224.3286-00.8709	   	 07 09 18.21 	   	 -10 31 52.88 	   	 21.39 	$\pm$	 0.09    	 II 	   	 d

60  	   	G224.3333-00.7476	   	 07 09 45.51 	   	 -10 28 43.81 	   	 33.59 	$\pm$	 0.13    	 I 	   	 e 	   	 ext. H$_2$ emission

61  	   	G224.3379-00.7384	   	 07 09 48.05 	   	 -10 28 42.87 	   	 29.73 	$\pm$	 0.31    	 II 	   	 d 	   	 near edge

62  	   	G224.3386-00.7425	   	 07 09 47.25 	   	 -10 28 51.88 	   	 08.02 	$\pm$	 0.27    	 II 	   	 d 	   	 near edge

63  	   	G224.3430-00.7437	   	 07 09 47.47 	   	 -10 29 08.11 	   	 18.24 	$\pm$	 0.30    	 II 	   	 d 	   	 near edge

64  	   	G224.3462-00.7394	   	 07 09 48.76 	   	 -10 29 10.99 	   	 5.18 	$\pm$	 0.12    	 I 	   	 e+d 	   	near edge, ext. H$_2$ emission

65  	   	G224.3684-00.7267	   	 07 09 53.94 	   	 -10 30 00.37 	   	 152.96 	$\pm$	 1.71	   	 II 	   	 e+d

66  	   	G224.3710-00.8882	   	 07 09 19.27 	   	 -10 34 36.14 	   	 07.37 	$\pm$	 0.14	   	 II 	   	 null 	   	 near edge

68  	   	G224.3955-00.7851	   	 07 09 44.35 	   	 -10 33 04.55 	   	 67.11 	$\pm$	 0.99	   	 II 	   	 d

69c   	   	G224.3986-00.7060	   	 07 10 01.84 	   	 -10 31 02.49 	   	 00.08 	$\pm$	 0.14	   	 I 	   	 null 	   	 integrated cont.

70  	   	G224.4042-00.7319	   	 07 09 56.84 	   	 -10 32 03.40 	   	 08.36 	$\pm$	 0.08	   	 II 	   	 d

71  	   	G224.4047-00.6971	   	 07 10 04.45 	   	 -10 31 07.18 	   	 16.38 	$\pm$	 0.26	   	 II 	   	 d

72  	   	G224.4055-00.6989	   	 07 10 04.13 	   	 -10 31 12.79 	   	 14.10 	$\pm$	 0.34	   	 II 	   	 null 	   	uncorr. H$_2$ emissione

73  	   	G224.4111-00.7059	   	 07 10 03.25 	   	 -10 31 42.93 	   	 10.08 	$\pm$	 0.04	   	 I 	   	 d

74  	   	G224.4114-00.6969	   	 07 10 05.25 	   	 -10 31 28.33 	   	 3.29 	$\pm$	 0.15	   	 II 	   	 null

75  	   	G224.4157-00.7023	   	 07 10 04.57 	   	 -10 31 50.47 	   	 6.83 	$\pm$	 0.15	   	 II 	   	 d

76  	   	G224.4195-00.7372	   	 07 09 57.43 	   	 -10 33 00.55 	   	 11.10 	$\pm$	 0.33	   	 II 	   	 d

77  	   	G224.4210-00.7045	   	 07 10 04.67 	   	 -10 32 11.44 	   	 2.83 	$\pm$	 0.03	   	 I 	   	 e+d

78  	   	G224.4238-00.6944	   	 07 10 07.16 	   	 -10 32 04.00 	   	 13.27 	$\pm$	 0.05	   	 II 	   	 null 	   	uncorr. H$_2$ emission

79  	   	G224.4259-00.6877	   	 07 10 08.87 	   	 -10 31 59.31 	   	 6.82 	$\pm$	 0.16	   	 II 	   	 d

80  	   	G224.4266-00.6996	   	 07 10 06.37 	   	 -10 32 21.11 	   	 35.91 	$\pm$	 0.22	   	 II 	   	 null 	   	uncorr. H$_2$ emission

81  	   	G224.4285-00.6662	   	 07 10 13.83 	   	 -10 31 31.86 	   	 25.17 	$\pm$	 0.75	   	 II 	   	 d

82A 	   	G224.4309-00.6943	   	 07 10 07.97 	   	 -10 32 26.94 	   	 08.97 	$\pm$	 0.11	   	 II 	   	 e+d

82B 	   	G224.4309-00.6943	   	 07 10 08.00 	   	 -10 32 26.34 	   	 07.96 	$\pm$	 0.25	   	  	   	 	   	 separation 663 au

83  	   	G224.4338-00.7003	   	 07 10 07.03 	   	 -10 32 45.39 	   	 4.11 	$\pm$	 0.21	   	 I 	   	 e+d 	   	 ext. H$_2$ emission?

84  	   	G224.4356-00.7049	   	 07 10 06.25 	   	 -10 32 59.03 	   	 67.02 	$\pm$	 1.28	   	 II 	   	 e

85  	   	G224.4361-00.6756	   	 07 10 12.64 	   	 -10 32 11.47 	   	 26.07 	$\pm$	 0.07	   	 II 	   	 null

86A 	   	G224.4404-00.7799	   	 07 09 50.55 	   	 -10 35 18.30 	   	 08.66 	$\pm$	 0.08	   	 II 	   	 d

86B 	   	G224.4404-00.7799	   	 07 09 50.55 	   	 -10 35 19.30 	   	 5.73 	$\pm$	 0.08	   	  	   	 	   	 separation 920 au

87A 	   	G224.4410-00.6702	   	 07 10 14.35 	   	 -10 32 18.18 	   	 1.89 	$\pm$	 0.11	   	 I 	   	 d 	   	 ext. H$_2$ emission?

87B 	   	G224.4410-00.6702	   	 07 10 14.42 	   	 -10 32 18.38 	   	 1.58 	$\pm$	 0.17	   	  	   	 	   	 separation 938 au

88	   	G224.4514-00.7054	   		   		   		   	 	   	 	   	not det., uncorr. H$_2$ emission

89A 	   	G224.4528-00.6930	   	 07 10 10.73 	   	 -10 33 34.32 	   	 4.21 	$\pm$	 0.03	   	 II 	   	 null

89B 	   	G224.4528-00.6930	   	 07 10 10.76 	   	 -10 33 34.32 	   	 4.27 	$\pm$	 0.05	     	   	 	   	hardly resolved, sep. 368 au

90  	   	G224.4564-00.6744	   	 07 10 15.20 	   	 -10 33 13.83 	   	 09.90 	$\pm$	 0.15    	 I 	   	 null

91A 	   	G224.4567-00.6781	   	 07 10 14.44 	   	 -10 33 22.29 	   	 91.33 	$\pm$	 1.19	   	 I 	   	 null 	   	 ext. H$_2$ emission

91B 	   	G224.4567-00.6781	   	 07 10 14.43 	   	 -10 33 21.49 	   	 22.22 	$\pm$	 0.10      	   	  	   	ext. H$_2$ emission, sep. 759 au

92  	   	G224.4577-00.7046	   	 07 10 08.80 	   	 -10 34 09.34 	   	 17.79 	$\pm$	 0.57    	 null 	   	 d

93  	   	G224.4578-00.6621	   	 07 10 18.00 	   	 -10 32 58.47 	   	 16.30 	$\pm$	 0.06    	 II 	   	 d

94  	   	G224.4593-00.6284	   	 07 10 25.47 	   	 -10 32 07.09 	   	 17.32 	$\pm$	 0.36    	 II 	   	 d

95  	   	G224.4599-00.6582	   	 07 10 19.06 	   	 -10 32 58.73 	   	 6.80 	$\pm$	 0.14    	 II 	   	 d

96  	   	G224.4649-00.6415	   	 07 10 23.25 	   	 -10 32 46.55 	   	 42.58 	$\pm$	 1.03    	 II 	   	 e

97  	   	G224.4664-00.6346	   	 07 10 24.93 	   	 -10 32 40.28 	   	 78.36 	$\pm$	 0.87	   	 II 	   	 d

98  	   	G224.4669-00.6489	   	 07 10 21.89 	   	 -10 33 05.39 	   	 14.29 	$\pm$	 0.16	   	 II 	   	 d

99  	   	G224.4677-00.6351	   	 07 10 24.98 	   	 -10 32 45.19 	   	 12.05 	$\pm$	 0.10	   	 I 	   	 null

100 	   	G224.4687-00.6368	   	 07 10 24.70 	   	 -10 32 51.26 	   	 12.85 	$\pm$	 0.37	   	 I 	   	 null 	   	 ext. H$_2$ emission

101 	   	G224.4697-00.6386	   	 07 10 24.48 	   	 -10 32 56.56 	   	 10.41 	$\pm$	 0.11	   	 II 	   	 null 	   	 near edge

102 	   	G224.4721-00.6296	   	 07 10 26.65 	   	 -10 32 50.22 	   	 94.01 	$\pm$	 0.93	   	 II 	   	 d

103 	   	G224.4746-00.6328	   	 07 10 26.27 	   	 -10 33 03.10 	   	 16.35 	$\pm$	 0.16	   	 I 	   	 null

104 	   	G224.4751-00.6383	   	 07 10 25.10 	   	 -10 33 14.00 	   	 54.74 	$\pm$	 0.42	   	 II 	   	 d

105 	   	G224.4793-00.6309	   	 07 10 27.20 	   	 -10 33 15.10 	   	 3.60 	$\pm$	 0.05	   	 II 	   	 d

106 	   	G224.4806-00.6430	   	 07 10 24.74 	   	 -10 33 39.50 	   	 1.94 	$\pm$	 0.23	   	 I 	   	 null

108 	   	G224.4838-00.7128	   	 07 10 09.97 	   	 -10 35 46.15 	   	 14.17 	$\pm$	 0.04	   	 II 	   	 d

109 	   	G224.4856-00.6229	   	 07 10 29.63 	   	 -10 33 22.03 	   	 08.77 	$\pm$	 0.20	   	 I 	   	 e+d 	   	 ext. H$_2$ emission

110 	   	G224.4859-00.6880	   	 07 10 15.56 	   	 -10 35 11.36 	   	 07.41 	$\pm$	 0.07	   	 II 	   	 d

111 	   	G224.4883-00.6555	   	 07 10 22.90 	   	 -10 34 24.70 	   	 33.94 	$\pm$	 0.12	   	 II 	   	 d 	   	ext. H$_2$ emission

112A 	   	G224.4908-00.6223	   	 07 10 30.41 	   	 -10 33 37.50 	   	 2.25 	$\pm$	 0.06	   	 I 	   	 d

112B 	   	G224.4908-00.6223	   	 07 10 30.36 	   	 -10 33 37.50 	   	 2.14 	$\pm$	 0.06    	  	   	 	   	 separation 736 au

113 	   	G224.4916-00.6754	   	 07 10 18.97 	   	 -10 35 08.47 	   	 07.11 	$\pm$	 0.13    	 II 	   	 d

114 	   	G224.4921-00.6620	   	 07 10 21.94 	   	 -10 34 47.82 	   	 9.08 	$\pm$	 0.04    	 II 	   	 d 	   	 ext. H$_2$ emission

115 	   	G224.4942-00.6660	   	 07 10 21.28 	   	 -10 35 00.91 	   	 19.05 	$\pm$	 0.21    	 II 	   	 e+d 	   	uncorr. H$_2$ emissione

116 	   	G224.4959-00.6715	   	 07 10 20.33 	   	 -10 35 15.87 	   	 23.58 	$\pm$	 0.15    	 II 	   	 d

117 	   	G224.4965-00.6805	   	 07 10 18.39 	   	 -10 35 32.83 	   	 47.41 	$\pm$	 0.58    	 II 	   	 d

118 	   	G224.5109-00.6777	   	 07 10 20.64 	   	 -10 36 14.23 	   	 13.32 	$\pm$	 0.09    	 II 	   	 d

aGLIMPSE360 IRAC Designations are SSTGLMA' followed by the names listed in this column; IRAC Designations are based on Galactic coordinates [Meade et al., 2014].
bThe YSO classification from [Sewiło et al., 2019]: the YSO Class and components (an envelope and/or a disk) identified based on the SED fitting with the [Robitaille, 2017] YSO models (e - envelope, e+d - envelope and disk, d - disk-only); null' indicates that the SED fitting results are not provided in [Sewiło et al., 2019].
Equatorial coordinates obtained from the integrated continuum maps in units of 10$^{-18}$ erg s$^{-1}$ cm$^{-2}$. See text for details.
dThe near edge label is given to sources located off-center (less than 0.8 from the map's edge).
eThe uncorr. H$_2$ emission  note indicates the presence of the extended H$_2$ emission in the field that is unrelated to the studied source.
Resolved binary candidates are marked with A and B for the brighter and fainter source, respectively. The separation between the two binary components is given in au in the Remarks for component B. The pixel size (0.2$\arcsec$) corresponds to 184 au.

The spectra of single sources were extracted using an aperture centered on the continuum peak with a radius of 3 pixels,  typically containing $\sim$70% of the total flux. For the binary candidates, different aperture sizes were used depending on the separation between the stars. The spectra were subsequently aperture-corrected using the ratio of the continuum emission within the aperture to that in the entire field of view; see Appendix <ref> for details.

Since KMOS is a moderate-resolution spectrograph, we assume a simplified shape of a spectral line and calculate integrated fluxes using Gaussian fits. The spectral line maps are constructed based on line fluxes estimated separately for each spaxel by fitting a Gaussian function to the emission line 50 times with randomly generated input parameters and choosing the fit with the lowest relative error.  All fits were visually confirmed. The same approach is applied to fitting the spectral lines in the extracted spectra.

KMOS $K$-band continuum maps of YSO candidates with possible stellar companions. Only spaxels with the continuum flux above 3$\sigma$ are shown. White contours correspond to 25%, 50%, and 75% of the $K$-band continuum peak. The position of the brighter source in each map is indicated with a red '$\times$' symbol (the A component in Table <ref>), and the fainter one with a grey '$+$' symbol (the B component). The size of the pixel (0.2$\arcsec$) corresponds to 184 au.

§ DATA ANALYSIS AND RESULTS

Out of 118 Spitzer YSO candidates observed with KMOS, 5 were not detected and 11 were resolved into two near-IR sources. In total, our KMOS observations provide the $K$-band data for 124 sources.

The IFU observations deliver both the spectral and spatial information. The $K$-band continuum emission allows us to identify all the continuum components and determine their positions and $K$-band fluxes. The spatial extent of the atomic and molecular emission hints at the underlying physical mechanisms and their characteristics. The hydrogen Br$\gamma$ emission is correlated with the UV continuum excess diagnostic of mass accretion. The CO ro-vibrational band head traces the inner disk. The H$_2$ emission lines, including the bright 1-0 S(1) line at 2.12 $\mu$m,  trace jets and shocks.

In this section, we discuss the spatial distribution of gas and dust, and summarize the line detections and kinematic information.

§.§ $K$-Band Continuum Emission

Example spectra of YSOs in CMa-$\ell$224. Source names from Table <ref> (No.) are provided on the left side of the respective spectra. All spectra are normalised to the continuum peak flux density in the wavelength range from 1.97 to 2.47 $\mu$m. The red horizontal lines show spectral ranges most affected by telluric lines; the wavelength range around 2.0 $\mu$m has been removed for clarity of the plot.

All the $K$-band continuum sources detected with VLT/KMOS in CMa-$\ell$224 are listed in Table <ref>. The table provides the sources' GLIMPSE360 catalog names and assigned KMOS IDs, their equatorial coordinates corresponding to the position of the $K$-band continuum peak, the $K$-band continuum fluxes estimated near 2.12 $\mu$m, the YSO classification from [Sewiło et al., 2019], and remarks on pointing, multiplicity, and the association with extended H$_2$ emission. Two $K$-band continuum components have been detected toward eleven fields centered on single Spitzer YSO candidates (see Fig. <ref>); the brighter source is listed in Table <ref> as source A” and the fainter one as source B.”

Selected $K$-band continuum maps are shown in Figure <ref> (11 fields with the detection of two sources) and Appendix <ref> (the remaining fields).

We consider a possibility that two sources detected in a single KMOS field are two components of a binary star/YSO. The median separation between the two components of the “binary star candidates” in CMa-$\ell$224 is 760 au (see Table <ref>), with a range from 370 au (source No. 89) to 1180 au (source No. 38).  All but three pairs of visual binary star candidates have spectra typical of accreting young stars: emission of Br$\gamma$ and/or H$_2$. Two more pairs have CO bandhead in absorption typical for cool stars. Given the proximity of the sources, they may be physically related. If they are binary systems, the multiplicity fraction defined as the ratio of the number of multiple systems to the total number of systems (both single and multiple), is $0.09\pm0.01$[We follow [Tobin et al., 2016] to estimate the uncertainty of derived multiplicity and adopt a bimodal statistic: $\sigma_{\rm{multi}}$=($N_{\rm{multi}}\cdot(1-N_{\rm{multi}}/N_{\rm{sys}})$)$^{-0.5}$ / $N_{\rm{sys}}$, where $N_{\rm{multi}}$ is a number of multiple systems and $N_{\rm{sys}}$ is a total number of systems.]. For comparison,  the multiplicity fraction for separations between 100 and 1000 au is $0.15\pm0.04$ for protostars in Perseus [Tobin et al., 2016], $0.14\pm0.01$ for protostars and $0.13\pm0.01$ for Pre-Main Sequence (PMS) stars in the Orion molecular clouds [Kounkel et al., 2016]. Adopting the approximate range recovered by KMOS (180-1180 au), the multiplicity fraction of protostars would be $\sim$0.09 in Perseus and 0.12 in Orion.

In summary, the binary star candidates in CMa-$\ell$224 have the multiplicity fraction that is consistent with YSOs in other molecular clouds, once we account for close binaries that are unresolved with KMOS. A more robust comparison requires an unbiased survey of YSOs at their earliest evolutionary stages and will be presented in the second paper in this series.

§.§ Spectral Line Detections

KMOS spectra contain several atomic and molecular lines emerging as a result of mass accretion and/or ejection in YSOs. Additionally, transitions of many atomic absorption lines which originate from a stellar photosphere are also detected and can be used to estimate spectral types and luminosity classes of YSOs <cit.>.

Figure <ref> shows the $K$-band spectra toward selected YSO candidates with diverse characteristics to illustrate the differences between the sources. For example, sources No. 47 and 19 show a strong H$_2$ emission, with only the latter showing also clear Br$\gamma$ and CO detections. Sources No. 12 and 25 both show CO bandhead in absorption. Yet, only source No. 12 is characterized also by a broad emission line profile in Br$\gamma$. Source No. 92 shows very weak H$_2$ emission, and Br$\gamma$ line in absorption. The line detections are discussed further in the subsequent sections.

§.§.§ Br$\gamma$

The most commonly detected line is the 2.1655 $\mu$m Br$\gamma$ hydrogen line, seen in emission in 58 sources (47%) and in absorption in 16 sources (13%). Twenty six sources with the Br$\gamma$ emission also show the CO lines in absorption. A similar pattern but higher detection rates of the Br$\gamma$ emission ($\sim75$%), was found toward a sample of 19 Class I protostars [Connelley & Greene, 2010, Connelley & Greene, 2014]. In a recent KMOS survey of YSOs in Perseus, the Br$\gamma$ was indeed more commonly detected in emission toward Class I (59%) than Class II YSOs <cit.>.

Three-color composite images of YSO candidates with the $\geq$3$\sigma$ detection in the H$_2$ line at 2.12 $\mu$m (in blue), Br$\gamma$ line at 2.17 $\mu$m (in green), and the $K$-band continuum (in red). In each image, white contours correspond to the $K$-band continuum emission with contour levels of 25%, 50%, and 75% of the continuum peak marked with the red '$\times$' symbol. Source IDs are indicated in the upper right corners.

The Br$\gamma$ line emission  in our sources is compact and associated with the continuum emission (Section <ref>, see also Fig. <ref> and Appendix <ref>). In low-mass YSOs, Br$\gamma$ emission originates from the close vicinity of the young star and usually traces magnetospheric accretion columns,
with possible contributions from winds [Kraus et al., 2008].
Some emission may also arise along the outflows [Beck et al., 2010].

The spatial extent of atomic emission of  Br$\gamma$-emitting gas is often more compact than molecular gas, with the exception of source No. 54, a bright star with the $K$-band magnitude of 9.9 (Fig. <ref>). Extended emission around a young star can be partially explained by an emerging H2 region, but a broad line profile suggests that other physical processes  might also contribute [Cooper et al., 2013]. The physical extent of $\sim$2350 au of Br$_\gamma$ emission from the source No. 54 cannot be explained with disk winds. In the subsequent analysis, the line fluxes are extracted from the area covering the continuum source, where
Br$\gamma$ emission is likely dominated by accretion.

§.§.§ H$_2$

H$_2$ is a homonuclear molecule that lacks a permanent dipole moment, thus only quadrupole transitions may occur ($\Delta J = 0, \pm2$). The extended distribution of the H$_2$ emission and its velocity structure indicate  the origin in bipolar outflows (Section <ref>). H$_2$ emission is detected in 33 sources in CMa-$\ell$224 (27%), including 14 that also show a detection of the Br$\gamma$ line in emission, and only 3 with the CO ro-vibrational lines in absorption. A similar H$_2$ detection rate of 23% was reported in a pioneering survey of Carr, 1990. Larger rates have been found toward Class I protostars <cit.> and higher-mass YSOs (56%, Cooper et al., 2013; 76%, Varricatt et al., 2010). Line fluxes are reported in Appendix <ref> in Table <ref>.

§.§.§ CO

The CO bandhead in absorption traces the atmosphere of a star or a disk. Among YSO candidates in CMa-$\ell$224, the CO first overtone is detected in absorption in 60% of YSO candidates, predominately (78%) in those classified as Class II <cit.>. Among Class I YSOs, the CO absorption is detected in 43% of the sources, similar to the 57% detection rate in [Connelley & Greene, 2010], and $>$2.5 times more than in the [Cooper et al., 2013] survey.

The H$_2$ 1-0 S(1) channel maps for source No. 47. The velocities relative to the rest velocity of the H$_2$ line are shown in the upper right corners. The green '$\times$` symbol shows the position of the $K$-band continuum peak.

The CO bandhead in emission indicates high accretion rates onto the YSO disk, possibly also associated with an accretion burst  <cit.>.
The CO bandhead emission requires high temperatures ($2500-4000$ K) and densities ($n_{\rm{H}}>10^{10}$ cm$^{-3}$) to excite the upper levels <cit.>. In CMa-$\ell$224, only 5% of the sources detected with KMOS show the CO bandhead in emission (4 Class I and 2 Class II YSOs). This detection rate is a factor of 4-5 lower than in the [Carr, 1989] and [Connelley & Greene, 2010] surveys, likely due their focus on Class I protostars. All of our sources with the CO bandhead emission also show the H$_2$ and Br$\gamma$ emission, and in four  out of six sources the H$_2$ emission is extended, suggesting intensive mass accretion and subsequent ejection.

§.§.§ Other Lines

Spectra of eleven YSO candidates show multiple atomic lines e.g., He I, Mg I, Ca I, Na I, K I, or Si I  typical for late spectral types  <cit.>; however, they do not reveal the presence of the Br$\gamma$, H$_2$, and CO lines. The confirmation of their YSO status would require careful spectral typing and further analysis, which will be presented in the future work in this series. Appendix <ref> shows spectra of all those sources and the information about the detection of the key spectral lines.

The H$_2$ 1-0 S(1) velocity centroid maps for all YSOs associated with the H$_2$ emission; the velocities are calculated with respect to the radial velocity of the corresponding source. Only pixels with the H$_2$ detection above 3$\sigma$ are shown. The contours show the $K$-band continuum emission, and the red crosses the positions of the YSO candidates (the peak of the $K$-band continuum emission). The source names are indicated in the upper right corners. Radial velocities of the gas associated with sources No. 78 and 88 cannot be estimated, and here only the apparent radial velocities are shown. The KMOS spectral sampling at 2.12 $\mu$m corresponds to 39.67 km/s.

§.§ Extinction

The line fluxes in the near-IR are affected by extinction. We adopt $A_{\rm{V}}$ obtained from the spectral energy distribution (SED) modelling by [Sewiło et al., 2019] for 85 sources in our sample. Our $K$-band observations cover several H$_2$ lines that are often used to estimate $A_{\rm{V}}$; however, after assessing the quality of our H$_2$ data, we decided against using this method due to the small separation between the two lines ($\sim0.2$ $\mu$m), uncertainty in telluric corrections near the Q-branch H$_2$ emission, and small number of sources for which this method can be used.

We assess the uncertainty of extinction estimated from the SED fitting using the work by [Furlan et al., 2016]. There, the R statistics was used to calculate a mode – the value with the highest frequency within a certain range from the best-fit R value (see their Fig. 48). Almost all their $A_{\rm{V}}$ uncertainties are consistent with a 1:1 correlation within a 10 mag error. We adopt this conservative value in the subsequent analysis, in particular in the calculation of the mass accretion
(Sections <ref>).

We note that extinction is expected to differ between the position of the infrared continuum source and the outflow [Yoon et al., 2022]. The ratios of the hydrogen Br$\gamma$ and Pa$\alpha$ lines are routinely used to estimate $A_{\rm{V}}$ toward the accretion region [Fiorellino et al., 2021], but the latter line is located in the $H$-band and not covered by our observations.
Since SED modelling provides the estimate of $A_{\rm{V}}$ toward the central star and H$_2$ emission is often extended, such $A_{\rm{V}}$ estimate is likely not representative for the true reddening of the H$_2$ lines. Therefore, we do not apply it for those lines. Thus, Table <ref> shows the H$_2$ line fluxes, which are not corrected for extinction.

§.§ Gas Spatial Distribution and Kinematics

The spatial distribution of the line emission varies significantly between different species detected with KMOS (Fig. <ref>). The Br$\gamma$ emission is spatially unresolved for the majority of the sources, in line with the origin in the inner disk. Similarly, the CO bandhead emission is also compact and co-spatial with the continuum emission. In contrast to the Br$\gamma$ and CO bandhead emission, the H$_2$ emission is  usually extended.

The extended distribution of the H$_2$ emission resembles the structure of the molecular outflows from Class 0/I protostars for 14 out of 33 sources with the H$_2$ line detection (42% or 11% of the entire sample; see Fig. <ref>). In five sources, the extended H$_2$ emission is not associated with the targeted source (see Table <ref>) and likely originates from an outflow from another YSO. Those five cases often show only one lobe within the field of view.
Among other sources with the  detection of H$_2$, the emission pattern does not show two prominent outflow lobes.
The lack of the bipolar pattern might be due to too low spatial resolution or the projection effects.
In general, a contribution from a photo-dissociation region is unlikely among the low-mass YSOs found in CMa-$\ell$224. The emission maps for the Br$\gamma$, H$_2$, and CO bandhead lines for the entire KMOS YSO sample in CMa-$\ell$224 are presented in Appendix <ref>.

We use the brightest and spectroscopically well-resolved H$_2$ line in $K$-band, the 1-0 S(1) transition at 2.12 $\mu$m, to calculate radial velocities of the emitting gas. Figure <ref> shows selected channel maps for an example source (No. 47) associated with an extended H$_2$ emission. Here, the blue-shifted emission transitions into red-shifted one on the other side of the central source, which is a clear signature of a bipolar ejection. The images reveal a velocity structure consistent with an outflow with velocities up to 50 km s$^{-1}$. A similar kinematic structure has been detected toward most of the other sources with the extended H$_2$ emission (No. 2, 15, 19, 46, 47, 54, 60, 64, 91 and 114; see Figure <ref>). The large spatial scales of those structures cannot be due to the emission from protoplanetary disks since they are too cold at such large distances from the central star to excite H$_2$.

The H$_2$ emission structure associated with sources No. 78 and 88 resembles the single outflow lobes; however, the continuum source is not seen in the map of the YSO candidate No. 88, while for No. 78, the H$_2$ emission does not seem to be physically connected to the source. Additional observations with a larger field-of-view are needed to identify sources responsible for the extended H$_2$ emission in these fields.

The estimated relative velocities do not exceed $\sim$50 km/s, with velocities increasing towards lobe ends, typical for outflow-driven shock waves. The moderate spectral resolution of KMOS does not allow for a detailed study on outflow kinematics, including any minor variations in kinematic patterns. The spatial extent of the H$_2$ gas supports the origin in outflow shocks. The limited wavelength coverage of our spectra, however, does not allow to properly estimate extinction as a function of position along the outflow (see Section <ref>).

Finally, we note that the bipolar structure in several H$_2$ lines toward source No. 47 (Appendix <ref>) is consistent with its classification as a Class I source using SED models [Sewiło et al., 2019]. The prominent H$_2$ outflow suggest that the source might be one of the youngest in our sample.
It would be a suitable candidate for a detailed spectral modeling using shock models; however, this is outside of the scope of this paper.

§.§ Gas Accretion

Gas accretion onto Class II YSOs is typically measured using UV continuum excess or line emission in e.g., H, He or Ca2 at optical wavelengths <cit.>. In more embedded Class I YSOs, the accretion is more often quantified using hydrogen lines in the near-IR which are less affected by dust extinction [Muzerolle et al., 1998, Natta et al., 2004].

We estimate the accretion luminosity ($L_\mathrm{acc}$ in units of L$_\odot$) for YSOs in CMa-$\ell$224 by measuring the flux of the hydrogen Br$\gamma$ line. As illustrated in Figure 3, the Br$\gamma$ emission arises from the same area as the continuum source; we therefore assume it traces primarily the unresolved accretion region (see also Section <ref>).
We use the relation between $L_{\mathrm{Br}\gamma}$ and $L_{\mathrm{acc}}$ for low- and intermediate-mass YSOs from  [Alcalá et al., 2017] to estimate $L_\mathrm{acc}$:

\begin{equation}
\label{eq:lacc}
\log{(L_{acc})} = (1.19\pm0.10)\cdot\log{L_{Br\gamma}} + (4.02\pm0.51)
\end{equation}

The mass accretion rates, $\dot{M}_{acc}$, are then calculated using the formula:

\begin{equation}
\label{eq:macc}
\dot{M}_{acc} = L_{acc} \times \left(1-\frac{R_*}{R_\mathrm{in}}\right)^{-1} \frac{R_*}{GM_*}
\end{equation}

where  $R_*$ and $M_*$ are the stellar radius and mass, respectively, and $R_{\mathrm{in}}$ is the inner-disk radius; typically $R_{\mathrm{in}}$ is assumed to equal 5$R_*$ [Gullbring et al., 1998]. Although the relation was found for Class II objects, it is a specific case of the general formula describing efficiency of converting kinetic energy of falling matter onto a stellar photosphere into radiation and can be applied to Class I sources as well (see the review of Hartmann et al., 2016 or a recent paper of Fiorellino et al., 2022).

Histogram of the accretion luminosities ($L_\mathrm{acc}$) based on Br$\gamma$ line luminosity (Table <ref>) measured toward YSOs in CMa-$\ell$224 (in dark orange). The median value of $L_\mathrm{acc}$ is indicated at the top of the plot. The orange and yellow  hatched histograms show the distribution of $L_\mathrm{acc}$ for sources classified by [Sewiło et al., 2019] as, respectively, Class I and Class II YSO candidates.

To calculate the mass accretion rates, knowledge of stellar masses and radii is needed. Stellar parameters for 17 sources from our sample have been estimated from the SED fitting[The spatial resolution of SED datapoints does not allow us to resolve visually multiple systems on the KMOS maps.
We therefore adopted the same stellar parameters for the two stars in a binary candidate, source  55 [Sewiło] et al., 2019]. Proper characterization of binaries is beyond the scope of this paper.
[Sewiło et al., 2019]; only half of them are associated with the Br$\gamma$ emission. The typical masses are $\sim2~M_\odot$ and radii are $\sim3~R_\odot$. The median logarithm of the accretion luminosity is -0.82$_{-0.84}^{+0.80}$
and the median logarithm of the accretion rate is -7.92$^{+1.06}_{-1.09}$
M$_\odot$yr$^{-1}$ for sources with $A_V$ estimates.

Accretion properties of YSOs in CMa-$\ell$224.

$\log{(L_{\rm{acc}})}$     $\log{(\dot{M}_{\mathrm{acc}})}$     $\log{(M_{\rm{out}})}$     $\log{(\dot{M}_{\rm{out}})}$

10$^{-16}$ (erg s$^{-1}$ cm$^{-2}$)     (L$_\odot$)     (M$_\odot$yr$^{-1}$)     (M$_\odot$)     (M$_\odot$yr$^{-1}$)

2     13.00 $\pm$ 0.32     -1.30$_{-0.82}^{+0.76}$         -6.07$_{-0.12}^{+0.07}$     -8.56$_{-0.15}^{+0.10}$

3     13.89 $\pm$ 0.32     -1.26$_{-0.73}^{+0.72}$

5     -397.29 $\pm$ 0.55

6     19.08 $\pm$ 0.25     -1.09$_{-0.81}^{+0.79}$

9     341.80 $\pm$ 0.79     0.40$_{-0.86}^{+0.86}$         -7.98$_{-0.17}^{+0.05}$     -10.15$_{-0.27}^{+0.13}$

10     54.74 $\pm$ 0.45     -0.55$_{-0.86}^{+0.84}$

12     352.30 $\pm$ 1.35     0.41$_{-0.90}^{+0.89}$

13     28.33 $\pm$ 0.77     -0.89$_{-0.80}^{+0.79}$         -7.45$_{-0.10}^{+0.03}$     -10.17$_{-0.32}^{+0.18}$

15     521.01 $\pm$ 0.35     0.62$_{-0.88}^{+0.87}$         -6.37$_{-0.15}^{+0.03}$     -8.42$_{-0.30}^{+0.14}$

16     108.77 $\pm$ 1.57     -0.19$_{-0.87}^{+0.85}$     -7.50$_{-1.07}^{+1.05}$

18*     1.52 $\pm$ 0.34     -2.40$_{-0.10}^{+0.08}$

19     68.62 $\pm$ 0.21     -0.43$_{-0.84}^{+0.83}$         -5.90$_{-0.14}^{+0.02}$     -8.32$_{-0.25}^{+0.11}$

20     12.93 $\pm$ 0.85     -1.29$_{-0.84}^{+0.80}$

22*     1.63 $\pm$ 0.56     -2.36$_{-0.19}^{+0.13}$

23*     4.91 $\pm$ 1.36     -1.79$_{-0.19}^{+0.15}$

24     100.01 $\pm$ 0.29     -0.24$_{-0.80}^{+0.80}$

26     18.52 $\pm$ 0.33     -1.11$_{-0.93}^{+0.85}$

28     33.00 $\pm$ 0.60     -0.81$_{-0.75}^{+0.74}$

30     1.00 $\pm$ 0.20     -2.62$_{-0.74}^{+0.69}$

31*     1.16 $\pm$ 0.16     -2.54$_{-0.03}^{+0.02}$

32     27.83 $\pm$ 3.14     -0.90$_{-0.86}^{+0.82}$         -9.32$_{-0.11}^{+0.07}$     -11.27$_{-0.19}^{+0.14}$

33     17.17 $\pm$ 0.31     -1.15$_{-0.82}^{+0.80}$         -8.57$_{-0.15}^{+0.06}$     -10.99$_{-0.38}^{+0.21}$

34*     4.82 $\pm$ 0.34     -1.80$_{-0.06}^{+0.06}$

35     -20.72 $\pm$ 0.54             -8.73$_{-0.14}^{+0.05}$     -10.99$_{-0.59}^{+0.26}$

36A*     65.58 $\pm$ 2.41     -0.46$_{-0.15}^{+0.15}$     -7.69$_{-0.35}^{+0.35}$

38A*     2.36 $\pm$ 0.35     -2.17$_{-0.07}^{+0.06}$

38B*     0.36 $\pm$ 0.09     -3.15$_{-0.06}^{+0.03}$

39     33.48 $\pm$ 3.62     -0.80$_{-0.95}^{+0.88}$     -8.10$_{-1.15}^{+1.08}$     -7.39$_{-0.11}^{+0.06}$     -9.66$_{-0.32}^{+0.20}$

40     -2119.07 $\pm$ 12.85

41*     10.11 $\pm$ 0.59     -1.42$_{-0.08}^{+0.08}$

43     -1.29 $\pm$ 0.27

45     4.54 $\pm$ 0.27     -1.83$_{-1.09}^{+0.87}$

46     307.57 $\pm$ 5.92     0.34$_{-0.87}^{+0.86}$     -6.97$_{-1.07}^{+1.06}$     -6.39$_{-0.12}^{+0.05}$     -8.60$_{-0.20}^{+0.12}$

47     373.10 $\pm$ 0.37     0.44$_{-0.89}^{+0.86}$         -4.52$_{-0.13}^{+0.03}$     -6.70$_{-0.17}^{+0.06}$

48     1.37 $\pm$ 0.33     -2.45$_{-0.80}^{+0.73}$

49     119.13 $\pm$ 0.60     -0.15$_{-0.90}^{+0.87}$         -7.09$_{-0.21}^{+0.09}$     -9.53$_{-0.55}^{+0.27}$

50     94.51 $\pm$ 0.52     -0.27$_{-0.78}^{+0.78}$

52A*     2.38 $\pm$ 0.45     -2.17$_{-0.10}^{+0.08}$

52B*     1.56 $\pm$ 0.50     -2.39$_{-0.17}^{+0.12}$

53     -107.15 $\pm$ 0.76

54*     493.43 $\pm$ 19.10     0.59$_{-0.24}^{+0.24}$         -6.32$_{-0.09}^{+0.08}$     -8.34$_{-0.18}^{+0.15}$

55A     9.30 $\pm$ 1.35     -1.46$_{-0.79}^{+0.76}$     -8.41$_{-0.99}^{+0.96}$

56     19.69 $\pm$ 0.40     -1.08$_{-0.86}^{+0.82}$

57*     3.87 $\pm$ 0.93     -1.92$_{-0.15}^{+0.12}$

58     7.93 $\pm$ 1.03     -1.55$_{-0.79}^{+0.76}$

60     1163.21 $\pm$ 1.90     1.03$_{-0.87}^{+0.85}$         -5.55$_{-0.10}^{+0.03}$     -8.06$_{-0.21}^{+0.11}$

61     -3.22 $\pm$ 0.45

63     -34.89 $\pm$ 0.53

64     413.58 $\pm$ 0.40     0.50$_{-0.71}^{+0.71}$         -7.40$_{-0.12}^{+0.08}$     -9.74$_{-0.24}^{+0.17}$

68     2.63 $\pm$ 0.39     -2.12$_{-0.83}^{+0.76}$

71     -69.23 $\pm$ 0.25             -6.28$_{-0.15}^{+0.03}$     -8.09$_{-0.25}^{+0.11}$

72*     1.66 $\pm$ 0.58     -2.36$_{-0.20}^{+0.13}$         -8.53$_{-0.09}^{+0.08}$     -11.22$_{-0.22}^{+0.18}$

73     44.53 $\pm$ 0.30     -0.66$_{-0.81}^{+0.80}$         -8.64$_{-0.15}^{+0.04}$     -11.33$_{-0.54}^{+0.24}$

75     -58.02 $\pm$ 0.27

76     30.09 $\pm$ 0.38     -0.86$_{-0.84}^{+0.81}$

77     276.96 $\pm$ 0.21     0.29$_{-0.85}^{+0.85}$         -8.43$_{-0.14}^{+0.02}$     -11.16$_{-0.56}^{+0.56}$

79     9.33 $\pm$ 0.31     -1.46$_{-0.77}^{+0.75}$     -8.48$_{-0.97}^{+0.95}$

80*     2.65 $\pm$ 0.73     -2.11$_{-0.16}^{+0.12}$         -8.68$_{-0.10}^{+0.08}$     -10.95$_{-0.21}^{+0.17}$

81     65.21 $\pm$ 0.82     -0.46$_{-0.90}^{+0.87}$     -7.74$_{-1.10}^{+1.07}$

83     625.84 $\pm$ 0.35     0.71$_{-0.84}^{+0.78}$         -8.38$_{-0.11}^{+0.05}$     -10.60$_{-0.25}^{+0.15}$

84     -109.85 $\pm$ 0.87

85*     8.34 $\pm$ 0.50     -1.52$_{-0.08}^{+0.07}$

86B     1.10 $\pm$ 0.14     -2.57$_{-0.72}^{+0.68}$

89A*     0.81 $\pm$ 0.23     -2.72$_{-0.12}^{+0.07}$

89B*     0.24 $\pm$ 0.08     -3.36$_{-0.10}^{+0.04}$

91A*     -6.12 $\pm$ 1.11

91B*     6.90 $\pm$ 0.50     -1.62$_{-0.08}^{+0.07}$         -7.25$_{-0.09}^{+0.08}$     -9.87$_{-0.53}^{+0.29}$

92     -14.75 $\pm$ 0.51

93     115.79 $\pm$ 0.52     -0.16$_{-0.84}^{+0.83}$

94     -3.35 $\pm$ 0.25

96     -56.89 $\pm$ 0.71

97     29.29 $\pm$ 1.65     -0.87$_{-0.89}^{+0.84}$

98     4.91 $\pm$ 0.42     -1.80$_{-0.75}^{+0.73}$

99*     1.42 $\pm$ 0.25     -2.44$_{-0.07}^{+0.05}$

102     273.94 $\pm$ 3.28     0.28$_{-0.85}^{+0.84}$

104     -6.90 $\pm$ 0.81

109     716.77 $\pm$ 0.24     0.78$_{-0.73}^{+0.73}$         -6.59$_{-0.13}^{+0.06}$     -9.20$_{-0.73}^{+0.30}$

111     51.99 $\pm$ 0.37     -0.58$_{-0.76}^{+0.76}$         -6.91$_{-0.33}^{+0.32}$     -9.39$_{-0.66}^{+0.50}$

112A     9.55 $\pm$ 0.16     -1.45$_{-0.72}^{+0.71}$

112B     30.86 $\pm$ 0.14     -0.84$_{-0.76}^{+0.76}$

114     187.72 $\pm$ 0.76     0.09$_{-1.08}^{+0.97}$     -7.06$_{-1.28}^{+1.17}$     -5.68$_{-0.20}^{+0.07}$     -7.69$_{-8.25}^{+8.25}$

116     -3.91 $\pm$ 0.28

117     -4.64 $\pm$ 0.51

118     9.01 $\pm$ 0.43     -1.48$_{-0.78}^{+0.76}$

Negative flux values indicate absorption. Measurements for targets with asterisk have not been corrected for extinction. Flux uncertainties do not cover the extinction uncertainty of 10 mag, which is included in uncertainties of accretion luminosity and mass accretion rates. Uncertainties of $\log{(L_{\mathrm{acc}})}$ and $\log{(\dot{M}_{\mathrm{acc}})}$ do not include the systematic distance uncertainty.

All accretion luminosities and, where available, mass accretion rates and associated uncertainties are listed in Table <ref>. The uncertainties of mass accretion rates
account for the uncertainties in
line fluxes, $L_{\rm Br\gamma}$–$L_{\rm acc}$ relation, extinction ($A_\mathrm{V}$ of 10 mag, corresponding to $\sim$1 mag in $K$-band), and 0.1 dex for both M$_*$ and R$_*$.
We do not include, however, the uncertainty of the distance to the region, which is systematic.
The adopted value of 0.92 kpc from [Sewiło et al., 2019] is consistent with those from stellar studies
<cit.>, and recent estimates using ALMA observations <cit.>.

We also cross-checked the distance to CMa-l224 with the results from the 3D dust maps using photometry from Pan-STARRS 1 and 2MASS, as well as the Gaia parallaxes[The interactive web interface: <http://argonaut.skymaps.info/>] [Green et al., 2019]. The median distance  estimated this way for several targets in our study yields 1.26 kpc. The value  exceeds our adopted distance of 0.92 kpc, with the difference that is larger than uncertainties of previous studies and the dispersion of individual distance values. Thus, we consider 0.34 kpc as the uncertainty of the distance; it corresponds to the uncertainty of 0.27 dex in accretion luminosity and mass accretion rate.

Distribution of accretion luminosity in the main filament in CMa-$\ell$224 (blue color scale). Class I sources are marked with hexagon and Class II sources with star symbols.  Black circles mark targets without $A_{\rm{V}}$ estimation. Marked with grey circles are three regions identified by [Sewiło et al., 2019].

Figure <ref> shows the distribution of accretion luminosities among Class I and II YSOs, which range from  -2.5 to 0.6 $L_{\odot}$ in logarithm.
The distributions of $L_{\mathrm{acc}}$ are similar for both Class I and II YSOs,
but the number of Class II YSOs in this sample is two times higher than Class I YSOs. We do not find any correlation between accretion luminosity and stellar mass, but the number of mass measurements is too small  and the mass measurements are too uncertain
to draw reliable conclusions. Accretion rate estimations are feasible only for 8 Class II YSOs (including one binary candidate), with available mass and radius estimates; their logarithmic values range from -8.4 to -7.0

§ DISCUSSION

§.§ Star Formation in CMa-$\ell$224

Star formation in Canis Majoris is closely linked with a supernova explosion $\sim$1.5 Myr ago [Herbst & Assousa, 1977, Comeron et al., 1998]. The expanding shell, coincident with CMa-$\ell$224, might have compressed the interstellar medium in a turbulent flow and triggered the filament formation [Fischer et al., 2016, Sewiło et al., 2019]. The main filament in the region is highly supercritical [Olmi et al., 2016], which leads to gravitational fragmentation and the formation of clumps and cores [André, 2017]. Overall, the filament is associated with $\sim$700 starless cores and $\sim$250 prestellar cores [Elia et al., 2013], and the largest  concentration
of Class I YSOs in the entire Canis Major star-forming region [Fischer et al., 2016, Sewiło et al., 2019].

Integral Field Spectroscopy in $K$-band at high-angular resolution provides unique insights into star formation processes on scales of individual YSOs in the main filament of CMa-$\ell$224. As shown in Section <ref>, the compact emission in the Br$\gamma$ line traces the on-going accretion onto the young stars, whereas the spatial distribution and kinematics of H$_2$ pin-points molecular outflows and mass ejection. The question remains, to what extent the mass accretion and ejections rates in YSOs are linked with the properties of the filament and its possible evolution.

Figure <ref> shows the distribution of accretion luminosity in the main filament in CMa-$\ell$224 (Section <ref>). YSOs are preferentially located within three clusters” associated with bright far-IR continuum emission. Following the nomenclature introduced in [Sewiło et al., 2019], cluster C is known to coincide also with C$^{18}$O 2-1 emission and the peak of $^{13}$CO 1-0, with the latter encompassing the entire filament [Olmi et al., 2016]. The $^{13}$CO and C$^{18}$O line observations revealed a velocity gradient along the main filament, which extends from cluster C to A [Olmi et al., 2016] and is likely responsible for the formation of YSOs at the ridge connecting clusters B and C. Another velocity gradient was identified from a secondary filament which is roughly perpendicular to the main one at the position of cluster B [Olmi et al., 2016]. [Sewiło et al., 2019] suggest that this smaller filament has provided molecular gas reservoir necessary to initiate the star formation for cluster B, as evidenced by the most evolved YSO population in this cluster.

To investigate this scenario, we calculate cumulative distributions of accretion luminosities, $L_{\rm{acc}}$,
for the three clusters (Figure <ref>). The distributions of accretion luminosity are similar for clusters A and B, whereas cluster C seems to be shifted towards lower values.
Observed similarity between the accretion distributions of clusters A and B is confirmed by the two-sample Kolmogorov–Smirnov (KS) test that yields statistic of  0.2 and $p$-value of 0.99.
The KS test for the cluster A and C, and cluster B and C, yield
statistic of 0.30–0.35 and $p$-value of 0.55-0.47, not allowing to statistically distinguish the two distributions.
The median accretion luminosities also gradually increase from Cluster C (-0.89 $\log{(L_\odot)}$), to B (-0.56 $\log{(L_\odot)}$) and A (-0.37 $\log{(L_\odot)}$). Higher accretion rate is expected from less evolved stars, assuming similar masses. Thus, the evolutionary stage of Clusters A and B could be earlier than those of Cluster C, contrary to earlier works [Sewiło et al., 2019].

However, the distributions can be affected by the number of targets within the Clusters, i.e. the small number of statistics.
Adopting the YSO classification from [Sewiło et al., 2019], Cluster C contains 5 Class I YSOs and 20 Class II YSOs for which we calculated accretion luminosities.
In Cluster A, we measured $L_{\rm{acc}}$ in 2 Class I and 6 Class II YSOs.
Similarly, in Cluster B – 3 Class I and 3 Class II YSOs have accretion estimates.
Thus, Cluster C contains significantly more YSOs with measured accretion luminosities (25) than the two other clusters combined (14). As a result, it has the most representative sample of sources, characterized by a broad range of accretion

Cumulative distributions of accretion luminosities
for clusters A (solid), B (dashed) and C (dotted line), as defined in [Sewiło et al., 2019]. Sources without $A_{\rm{V}}$ estimates are excluded.

§.§ A Possible Impact of Low Metallicity

Observations of star formation in external galaxies suggest that subsolar metallicity has a clear impact on the physical and chemical conditions in molecular clouds <cit.>. In particular, a lower dust content in the low-metallicity environments reduces the shielding from ultraviolet radiation and translates to a decrease of molecular abundances due to photodissociation <cit.>. The formation rate of molecules on the dust grains also decreases, which is the main reason for the reduction of H$_2$ abundances in the low-metallicity clouds [Glover & Clark, 2012].

The Large and Small Magellanic Clouds, with Z of $\sim$0.3-0.5 Z$_{\odot}$ <cit.> and 0.2 Z$_{\odot}$ <cit.>, are routinely used as laboratories for studying star formation in low-metallicity environments. Observations show a decrease of molecular gas cooling in YSO envelopes, UV radiation penetrating deeper into the clouds, and an increase in average dust temperatures with respect to those in local star forming regions <cit.>. However,
the distances of 50.0$\pm$1.1 kpc (LMC, Pietrzyński et al., 2013) and 62.1$\pm$2.0 kpc (SMC, Graczyk et al., 2014) hinder characterisation of individual low-mass YSOs.

Accretion luminosity versus absolute $K$-band magnitude. Red stars show YSOs in CMa-$\ell$224, blue circles show MYSOs [Cooper et al., 2013], light-green +” symbols show MYSOs from SMC [Ward et al., 2017], green crosses show MYSOs from LMC [Ward et al., 2016], violet open squares show YSOs in Orion [Manara et al., 2012], and orange triangles show low-mass YSOs in Perseus [Fiorellino et al., 2021]. The dashed lines show linear fits separately to each of the datasets with the same colors. The grey line show a linear fit to all datapoints. All luminosities have been corrected for extinction.

Accretion luminosity versus star luminosity. The symbols are the same as in Fig. <ref>. The dashed line show a linear fit to the data. All luminosities have been corrected for extinction.

The outer Galaxy also consists of star forming regions with the subsolar metallicities, and so offers the opportunity to study the impact of low metal content on star formation [Sodroski et al., 1997]. The metallicity of CMa-$\ell$224 has not been measured, so here we adopt the O/H Galactocentric radial gradients obtained from the observations of HII regions  accross the Milky Way. Depending on the adopted survey, the metallicity of CMa-$\ell$224 would be 0.55-0.58 Z$_{\odot}$ [Esteban & García-Rojas, 2018], 0.68 Z$_{\odot}$ [Fernández-Martín et al., 2017], or 0.67-0.73 Z$_{\odot}$ [Balser et al., 2011]. The O/H and Fe/H gradients obtained using the observations of Cepheids suggest much higher metallicities of 1.11-1.17 Z$_{\odot}$ [Maciel & Andrievsky, 2019] or 1.13-1.27 Z$_{\odot}$ [Luck & Lambert, 2011], yet they are not representative for star forming regions. Therefore, we assume that the metallicity of CMa-$\ell$224 is intermediate between the metallicity of the LMC and the metallicity in the Solar neighborhood.

A possible effect of metallicity on YSO accretion has been investigated in the LMC and SMC based on the near-IR observations of massive YSOs [Ward et al., 2016, Ward et al., 2017], and optical observations with H$\alpha$ in YSOs with a wide range of masses [De Marchi et al., 2011, De Marchi et al., 2017, Biazzo et al., 2019]. The Br$\gamma$ luminosities for Magellanic YSOs seem higher than those found in the Galactic high-mass YSOs from the RMS survey for the same absolute magnitudes [Cooper et al., 2013]. Here, we extend the measurements toward lower-mass YSOs, which are located in CMa-$\ell$224 and more nearby star-forming regions (Figs. <ref> and <ref>). We adopt the same empirical relations between the Br$\gamma$ luminosity and accretion luminosities for various samples [Alcalá et al., 2017]; however, $L_{\mathrm{acc}}$ from Orion has been calculated using H$\alpha$ and two-color diagrams [Manara et al., 2012]. The $K$-band magnitudes of YSOs from Orion are adopted from [Robberto et al., 2010].

Figures <ref> and <ref> show the increase of $L_\mathrm{acc}$ with both the absolute $K$-band magnitude and $L_*$, with the YSOs from CMa-$\ell$224 located between the least massive YSOs in the NGC 1333 cluster in Perseus [Fiorellino et al., 2021] and in Orion [Manara et al., 2012], and the high-mass YSOs both in the Milky Way [Cooper et al., 2013] and the Magellanic Clouds [Ward et al., 2016, Ward et al., 2017]. The Pearson coefficient of $|r|=0.96$ (30$\sigma$) is obtained for the absolute $K$-band magnitude vs. $L_\mathrm{acc}$ relation,  reflecting a strong correlation (Fig. <ref>). Similarly, the Pearson coefficient of $|r|=0.93$ (30$\sigma$) characterises the $L_*$ vs $L_\mathrm{acc}$ relation.

Accretion luminosities in CMa-$\ell$224 are consistent with those in the Solar neighbourhood (e.g., Perseus, Orion) and do not show a clear enhancement of accretion luminosity expected in the low-metallicity environments [Ward et al., 2016, Ward et al., 2017]. Most of the Br$\gamma$ emitters are however Class II sources, characterized by lower $L_{\mathrm{acc}}$ with respect to less evolved, Class I sources.
The metallicity of CMa-$\ell$224 <cit.> might be too similar to the Solar neighbourhood to exhibit a clear, measurable difference in accretion properties.
Follow-up observations of additional clouds in the outer Galaxies, spanning a broad range of metallicities and sufficient number of sources, would be critical to draw robust conclusions about the impact of the lower metallicity on accretion rates in the Outer Galaxy.

§ CONCLUSIONS

We presented VLT/KMOS near-IR spectroscopy of YSOs in CMa-$\ell$224, illustrating the instrument capability to study low-mass star formation, and where applicable, to detect molecular outflows in the outer Galaxy. We summarise the findings as follows:

* The Br$\gamma$ line is detected in emission toward 47% and in absorption toward 13% of the sources.
The H$_2$ 1-0 S(1) line is detected on-source in 33 sources, including 27 sources where the emission is extended. The CO bandhead at 2.3 $\mu$m is detected in emission in 5% of YSOs and in absorption in 60% of YSOs.

* The extent and velocity structures of H$_2$ emission suggest an origin of the emission in outflows and shocks. Follow-up observations are necessary, however, to estimate the extinction along the outflows.

* We find a gradual increase in accretion luminosities from Cluster C to A and B with median values of -0.89, -0.56, and -0.37 $\log{L_\odot}$, respectively. Cluster C might be the most evolved part of CMa-$\ell$224.

* Accretion luminosities
do not show the impact of sub-Solar metallicity in the CMa-$\ell$224.
It is likely due to insufficient difference in metallicity between the region and
the Solar neighbourhood.

Large-scale galactic surveys, such as the Outer Galaxy High Resolution Survey [Colombo et al., 2021], are starting to provide a complete census of molecular clouds and star-forming filaments spanning a range of metallicities. IFU observations of YSOs in those regions, either with ground-based facilities or the James Webb Space Telescope, will be critical to confirm the impact of metallicity on low-mass star formation within our Galaxy.

The authors thank the referee for a careful reading of the manuscript.
DI, AK, MK, and NL acknowledge support from the First TEAM grant of the Foundation for Polish Science No. POIR.04.04.00-00-5D21/18-00. This work was supported by the Polish National Science Center grants 2014/15/B/ST9/02111 and 2016/21/D/ST9/01098. The material is based upon work supported by NASA under award number 80GSFC21M0002 (MS). This article has been supported by the Polish National Agency for Academic Exchange under Grant No. PPI/APM/2018/1/00036/U/001. DI was partly funded by the European Research Council (ERC) via the ERC Synergy Grant ECOGAL (grant 855130).
The research of LEK is supported by a research grant (19127) from VILLUM FONDEN. GJH is supported by general grant 12173003 awarded by the National Natural Science Foundation of China.
WRMR thanks the financial support from the Leiden Observatory.
Based on observations made with ESO Telescopes at the La Silla Paranal Observatory under programme ID 0102.C-0914(A).

astropy[<http://www.astropy.org>] [Astropy Collaboration et al., 2013, Price-Whelan et al., 2018],
matplotlib[<https://matplotlib.org/>] [Hunter, 2007],
numpy[<https://numpy.org/>] [Oliphant, 2006, Van Der Walt et al., 2011],

[Agra-Amboage et al., 2014]
Agra-Amboage, V., Cabrit, S., Dougados, C., et al. 2014, , 564,
A11, 10.1051/0004-6361/201220488

[Alcalá et al., 2014]
Alcalá, J. M., Natta, A., Manara, C. F., et al. 2014, , 561,
A2, 10.1051/0004-6361/201322254

[Alcalá et al., 2017]
Alcalá, J. M., Manara, C. F., Natta, A., et al. 2017, , 600,
A20, 10.1051/0004-6361/201629929

[André, 2017]
André, P. 2017, Comptes Rendus Geoscience, 349, 187,

[Ansdell et al., 2017]
Ansdell, M., Williams, J. P., Manara, C. F., et al. 2017, , 153,
240, 10.3847/1538-3881/aa69c0

[Astropy Collaboration et al., 2013]
Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013,
, 558, A33, 10.1051/0004-6361/201322068

[Bally, 2016]
Bally, J. 2016, , 54, 491, 10.1146/annurev-astro-081915-023341

[Balser et al., 2011]
Balser, D. S., Rood, R. T., Bania, T. M., & Anderson, L. D. 2011,
, 738, 27, 10.1088/0004-637X/738/1/27

[Beck et al., 2010]
Beck, T. L., Bary, J. S., & McGregor, P. J. 2010, , 722, 1360,

[Benedettini et al., 2020]
Benedettini, M., Molinari, S., Baldeschi, A., et al. 2020, , 633,
A147, 10.1051/0004-6361/201936096

[Biazzo et al., 2019]
Biazzo, K., Beccari, G., De Marchi, G., & Panagia, N. 2019, , 875,
51, 10.3847/1538-4357/ab0f95

[Bloemen et al., 1984]
Bloemen, J. B. G. M., Caraveo, P. A., Hermsen, W., et al. 1984, ,
139, 37

[Carr, 1989]
Carr, J. S. 1989, , 345, 522, 10.1086/167927

[Carr, 1990]
—. 1990, , 100, 1244, 10.1086/115591

[Casali & Eiroa, 1996]
Casali, M. M., & Eiroa, C. 1996, , 306, 427

[Claria, 1974]
Claria, J. J. 1974, , 79, 1022, 10.1086/111648

[Coccato et al., 2019]
Coccato, L., Freudling, W., Smette, A., et al. 2019, The Messenger,
177, 14, 10.18727/0722-6691/5148

[Colombo et al., 2021]
Colombo, D., König, C., Urquhart, J. S., et al. 2021, , 655,
L2, 10.1051/0004-6361/202142182

[Comeron et al., 1998]
Comeron, F., Torra, J., & Gomez, A. E. 1998, , 330, 975

[Connelley & Greene, 2010]
Connelley, M. S., & Greene, T. P. 2010, , 140, 1214,

[Connelley & Greene, 2014]
—. 2014, , 147, 125, 10.1088/0004-6256/147/6/125

[Cooper et al., 2013]
Cooper, H. D. B., Lumsden, S. L., Oudmaijer, R. D., et al. 2013,
, 430, 1125, 10.1093/mnras/sts681

[Cyganowski et al., 2011]
Cyganowski, C. J., Brogan, C. L., Hunter, T. R., Churchwell, E., &
Zhang, Q. 2011, , 729, 124, 10.1088/0004-637X/729/2/124

[Cyganowski et al., 2008]
Cyganowski, C. J., Whitney, B. A., Holden, E., et al. 2008, , 136,
2391, 10.1088/0004-6256/136/6/2391

[Davies et al., 2013]
Davies, R. I., Agudo Berbel, A., Wiezorrek, E., et al. 2013, , 558,
A56, 10.1051/0004-6361/201322282

[Davis & Eisloeffel, 1995]
Davis, C. J., & Eisloeffel, J. 1995, , 300, 851

[De Marchi et al., 2017]
De Marchi, G., Panagia, N., & Beccari, G. 2017, , 846, 110,

[De Marchi et al., 2011]
De Marchi, G., Panagia, N., Romaniello, M., et al. 2011, , 740, 11,

[Djordjevic et al., 2019]
Djordjevic, J. O., Thompson, M. A., Urquhart, J. S., & Forbrich, J.
2019, , 487, 1057, 10.1093/mnras/stz1262

[Elia et al., 2013]
Elia, D., Molinari, S., Fukui, Y., et al. 2013, , 772, 45,

[Esteban & García-Rojas, 2018]
Esteban, C., & García-Rojas, J. 2018, , 478, 2315,

[Fernández-Martín et al., 2017]
Fernández-Martín, A., Pérez-Montero, E., Vílchez,
J. M., & Mampaso, A. 2017, , 597, A84,

[Fiorellino et al., 2022]
Fiorellino, E., Tychoniec, Ł., Manara, C. F., et al. 2022, ,
937, L9, 10.3847/2041-8213/ac8fee

[Fiorellino et al., 2021]
Fiorellino, E., Manara, C. F., Nisini, B., et al. 2021, , 650, A43,

[Fischer et al., 2016]
Fischer, W. J., Padgett, D. L., Stapelfeldt, K. L., & Sewiło, M.
2016, , 827, 96, 10.3847/0004-637X/827/2/96

[Frank et al., 2014]
Frank, A., Ray, T. P., Cabrit, S., et al. 2014, in Protostars and
Planets VI, ed. H. Beuther, R. S. Klessen, C. P. Dullemond, &
T. Henning, 451, 10.2458/azu_uapress_9780816531240-ch020

[Furlan et al., 2016]
Furlan, E., Fischer, W. J., Ali, B., et al. 2016, , 224, 5,

[Glover & Clark, 2012]
Glover, S. C. O., & Clark, P. C. 2012, , 426, 377,

[Graczyk et al., 2014]
Graczyk, D., Pietrzyński, G., Thompson, I. B., et al. 2014, ,
780, 59, 10.1088/0004-637X/780/1/59

[Green et al., 2019]
Green, G. M., Schlafly, E., Zucker, C., Speagle, J. S., &
Finkbeiner, D. 2019, , 887, 93, 10.3847/1538-4357/ab5362

[Gullbring et al., 1998]
Gullbring, E., Hartmann, L., Briceño, C., & Calvet, N. 1998, ,
492, 323, 10.1086/305032

[Guo et al., 2021]
Guo, Z., Lucas, P. W., Contreras Peña, C., et al. 2021, ,
504, 830, 10.1093/mnras/stab882

[Hartmann et al., 2016]
Hartmann, L., Herczeg, G., & Calvet, N. 2016, , 54, 135,

[Hawkins, 2022]
Hawkins, K. 2022, arXiv e-prints, arXiv:2207.04542.

[Herbst & Assousa, 1977]
Herbst, W., & Assousa, G. E. 1977, , 217, 473, 10.1086/155596

[Heyer & Dame, 2015]
Heyer, M., & Dame, T. M. 2015, , 53, 583,

[Hunter, 2007]
Hunter, J. D. 2007, Computing In Science & Engineering, 9, 90

[Kaltcheva & Hilditch, 2000]
Kaltcheva, N. T., & Hilditch, R. W. 2000, , 312, 753,

[Karska et al., 2018]
Karska, A., Kaufman, M. J., Kristensen, L. E., et al. 2018, , 235,
30, 10.3847/1538-4365/aaaec5

[Kennicutt & Evans, 2012]
Kennicutt, R. C., & Evans, N. J. 2012, , 50, 531,

[Kounkel et al., 2016]
Kounkel, M., Megeath, S. T., Poteet, C. A., Fischer, W. J., &
Hartmann, L. 2016, , 821, 52, 10.3847/0004-637X/821/1/52

[Kraus et al., 2008]
Kraus, S., Hofmann, K. H., Benisty, M., et al. 2008, , 489, 1157,

[Kristensen et al., 2007]
Kristensen, L. E., Ravkilde, T. L., Field, D., Lemaire, J. L., &
Pineau Des Forêts, G. 2007, , 469, 561,

[Kristensen et al., 2017]
Kristensen, L. E., van Dishoeck, E. F., Mottram, J. C., et al. 2017,
, 605, A93, 10.1051/0004-6361/201630127

[Lépine et al., 2011]
Lépine, J. R. D., Cruz, P., Scarano, S., J., et al. 2011, ,
417, 698, 10.1111/j.1365-2966.2011.19314.x

[Lin et al., 2021]
Lin, Z., Sun, Y., Xu, Y., Yang, J., & Li, Y. 2021, , 252, 20,

[Lombardi et al., 2011]
Lombardi, M., Alves, J., & Lada, C. J. 2011, , 535, A16,

[Lorenzetti et al., 2009]
Lorenzetti, D., Larionov, V. M., Giannini, T., et al. 2009, , 693,
1056, 10.1088/0004-637X/693/2/1056

[Luck & Lambert, 2011]
Luck, R. E., & Lambert, D. L. 2011, , 142, 136,

[Luhman & Rieke, 1998]
Luhman, K. L., & Rieke, G. H. 1998, , 497, 354, 10.1086/305447

[Maciel & Andrievsky, 2019]
Maciel, W. J., & Andrievsky, S. 2019, arXiv e-prints, arXiv:1906.01686.

[Madden et al., 2013]
Madden, S. C., Rémy-Ruyer, A., Galametz, M., et al. 2013, ,
125, 600, 10.1086/671138

[Mainzer et al., 2011]
Mainzer, A., Grav, T., Bauer, J., et al. 2011, , 743, 156,

[Manara et al., 2012]
Manara, C. F., Robberto, M., Da Rio, N., et al. 2012, , 755, 154,

[Meade et al., 2014]
Meade, M., Whitney, B., & Babler, B. 2014, (Pasadena, CA: IPAC),

[Molinari et al., 2010]
Molinari, S., Swinyard, B., Bally, J., et al. 2010, , 518, L100,

[Mottram et al., 2017]
Mottram, J. C., van Dishoeck, E. F., Kristensen, L. E., et al. 2017,
, 600, A99, 10.1051/0004-6361/201628682

[Muzerolle et al., 1998]
Muzerolle, J., Hartmann, L., & Calvet, N. 1998, , 116, 2965,

[Natta et al., 2004]
Natta, A., Testi, L., Muzerolle, J., et al. 2004, , 424, 603,

[Nisini et al., 2018]
Nisini, B., Antoniucci, S., Alcalá, J. M., et al. 2018, , 609,
A87, 10.1051/0004-6361/201730834

[Nisini et al., 2005]
Nisini, B., Antoniucci, S., Giannini, T., & Lorenzetti, D. 2005, ,
429, 543, 10.1051/0004-6361:20041409

[Nisini et al., 2015]
Nisini, B., Santangelo, G., Giannini, T., et al. 2015, , 801, 121,

[Offner et al., 2011]
Offner, S. S. R., Lee, E. J., Goodman, A. A., & Arce, H. 2011, ,
743, 91, 10.1088/0004-637X/743/1/91

[Oliphant, 2006]
Oliphant, T. E. 2006, A guide to NumPy, Vol. 1 (Trelgol Publishing USA)

[Oliveira et al., 2019]
Oliveira, J. M., van Loon, J. T., Sewiło, M., et al. 2019, ,
490, 3909, 10.1093/mnras/stz2810

[Olmi et al., 2023]
Olmi, L., Brand, J., & Elia, D. 2023, , 518, 1917,

[Olmi et al., 2016]
Olmi, L., Cunningham, M., Elia, D., & Jones, P. 2016, , 594, A58,

[Pietrzyński et al., 2013]
Pietrzyński, G., Graczyk, D., Gieren, W., et al. 2013, , 495,
76, 10.1038/nature11878

[Predehl & Schmitt, 1995]
Predehl, P., & Schmitt, J. H. M. M. 1995, , 293, 889

[Price-Whelan et al., 2018]
Price-Whelan, A. M., Sipőcz, B. M., Günther, H. M., et al.
2018, , 156, 123, 10.3847/1538-3881/aabc4f

[Robberto et al., 2010]
Robberto, M., Soderblom, D. R., Scandariato, G., et al. 2010, , 139,
950, 10.1088/0004-6256/139/3/950

[Robitaille, 2017]
Robitaille, T. P. 2017a, , 600, A11,

[Robitaille, 2017]
—. 2017b, , 600, A11, 10.1051/0004-6361/201425486

[Rolleston et al., 2002]
Rolleston, W. R. J., Trundle, C., & Dufton, P. L. 2002, , 396, 53,

[Roman-Duval et al., 2010]
Roman-Duval, J., Jackson, J. M., Heyer, M., Rathborne, J., & Simon,
R. 2010, , 723, 492, 10.1088/0004-637X/723/1/492

[Roman-Duval et al., 2014]
Roman-Duval, J., Gordon, K. D., Meixner, M., et al. 2014, , 797,
86, 10.1088/0004-637X/797/2/86

[Russell & Dopita, 1992]
Russell, S. C., & Dopita, M. A. 1992, , 384, 508,

[Schisano et al., 2014]
Schisano, E., Rygl, K. L. J., Molinari, S., et al. 2014, , 791, 27,

[Sewiło et al., 2019]
Sewiło, M., Whitney, B. A., Yung, B. H. K., et al. 2019, , 240,
26, 10.3847/1538-4365/aaf86f

[Sharples et al., 2013]
Sharples, R., Bender, R., Agudo Berbel, A., et al. 2013, The Messenger,
151, 21

[Shevchenko et al., 1999]
Shevchenko, V. S., Ezhkova, O. V., Ibrahimov, M. A., van den Ancker,
M. E., & Tjin A Djie, H. R. E. 1999, , 310, 210,

[Skrutskie et al., 2006]
Skrutskie, M. F., Cutri, R. M., Stiening, R., et al. 2006, , 131,
1163, 10.1086/498708

[Sodroski et al., 1997]
Sodroski, T. J., Odegard, N., Arendt, R. G., et al. 1997, , 480,
173, 10.1086/303961

[Stanke et al., 2002]
Stanke, T., McCaughrean, M. J., & Zinnecker, H. 2002, , 392, 239,

[Testi et al., 2022]
Testi, L., Natta, A., Manara, C. F., et al. 2022, , 663, A98,

[Tobin et al., 2016]
Tobin, J. J., Looney, L. W., Li, Z.-Y., et al. 2016, The Astrophysical
Journal, 818, 73, 10.3847/0004-637x/818/1/73

[van der Marel et al., 2013]
van der Marel, N., Kristensen, L. E., Visser, R., et al. 2013, ,
556, A76, 10.1051/0004-6361/201220717

[Van Der Walt et al., 2011]
Van Der Walt, S., Colbert, S. C., & Varoquaux, G. 2011, Computing in Science
& Engineering, 13, 22

[van Kempen et al., 2010]
van Kempen, T. A., Kristensen, L. E., Herczeg, G. J., et al. 2010,
, 518, L121, 10.1051/0004-6361/201014615

[van Loon et al., 2010]
van Loon, J. T., Oliveira, J. M., Gordon, K. D., Sloan, G. C., &
Engelbracht, C. W. 2010a, , 139, 1553,

[van Loon et al., 2010]
van Loon, J. T., Oliveira, J. M., Gordon, K. D., et al.
2010b, , 139, 68, 10.1088/0004-6256/139/1/68

[Varricatt et al., 2010]
Varricatt, W. P., Davis, C. J., Ramsay, S., & Todd, S. P. 2010,
, 404, 661, 10.1111/j.1365-2966.2010.16356.x

[Visser et al., 2012]
Visser, R., Kristensen, L. E., Bruderer, S., et al. 2012, , 537,
A55, 10.1051/0004-6361/201117109

[Ward et al., 2016]
Ward, J. L., Oliveira, J. M., van Loon, J. T., & Sewiło, M. 2016,
, 455, 2345, 10.1093/mnras/stv2424

[Ward et al., 2017]
—. 2017, , 464, 1512, 10.1093/mnras/stw2386

[Westerlund, 1997]
Westerlund, B. E. 1997, The Magellanic Clouds

[Wolfire et al., 2010]
Wolfire, M. G., Hollenbach, D., & McKee, C. F. 2010, , 716, 1191,

[Wright et al., 2010]
Wright, E. L., Eisenhardt, P. R. M., Mainzer, A. K., et al. 2010, ,
140, 1868, 10.1088/0004-6256/140/6/1868

[Yoon et al., 2022]
Yoon, S.-Y., Herczeg, G. J., Lee, J.-E., et al. 2022, , 929, 60,

§ KMOS $K$-BAND CONTINUUM MAPS

Figure <ref> shows the $K$-band continuum flux density maps for YSO candidates in CMa-$\ell$224 without stellar companions. Similar maps for the fields with the detection of two sources (the binary star candidates) are shown in Figure <ref>.  Figure <ref> shows the integrated continuum maps for YSO candidates with the very weak $K$-band continuum.

KMOS $K$-band continuum maps for single YSO candidates. Source names are provided in the upper right corners (No. in Table <ref>). The contour levels are 0.25, 0.5, and 0.75$\times$ the  $K$-band continuum peak. Only pixels with the continuum flux density above 3$\sigma$ are shown.

Similar to Figure <ref>, but for the faintest sources in the sample: No. 11, 14, 26, and 69 (see Table <ref>). The continuum was integrated over the entire $K$-band spectrum in every pixel to increase the signal-to-noise ratio.

§ SPECTRA

Figures <ref>-<ref> show the $K$-band spectra of all YSO candidates in CMa-$\ell$224 observed with KMOS. The spectra are presented in subgroups having similar characteristics – line detections and their profiles. In particular,
Figure <ref> shows spectra of sources with the Br$\gamma$ line in emission, as represented by source No. 12 in Figure <ref>.
Figure <ref> presents spectra of sources with the detections of both the Br$\gamma$ and H$_2$ at 2.12 $\mu$m lines in emission, but a lack of the CO bandhead emission.
Figure <ref> shows spectra of two sources with the strong H$_2$ emission from outflows, but a lack of the Br$\gamma$ or CO bandhead lines.
Figure <ref> shows spectra of three sources with the Br$\gamma$ and H$_2$ lines in emission, and the CO bandhead lines in absorption.
Figure <ref> features spectra of presumably more evolved objects without the H$_2$ emission lines, but with the Br$\gamma$ line in emission and the CO bandhead lines in absorption.
Figure <ref> shows spectra of sources with the H$_2$ emission and the CO bandhead absorption lines, but a lack of the Br$\gamma$ emission.
Figure <ref> shows spectra of sources with the detection of Br$\gamma$ in absorption, and Figure <ref>, the only five single sources with spectra characterised by the CO bandhead in emission. Figure <ref> shows spectra of sources with CO in absorption, but non-detections of the H$_2$ or Br$\gamma$ lines. Finally, Figure <ref> shows spectra of YSO candidates which do not show any signatures of accretion or ejection, and require follow-up observations to confirm their status.

Spectra of the individual components of the binary YSO candidates are presented in Figures <ref>-<ref>. While all spectra of single objects were extracted from the spectral cubes within the aperture of radius of 3 pixels, due to their proximity, for double stars different approach has been followed. For the most distant pairs (sources No. 21, 38, 38, 86, and 87) the same aperture radius was used. For sources No. 82, 91, and 112 we used the radius of 2 pixels. To maximise the collected signal, we used the radius of 1.5 pixels for pairs No. 52 and 55. Spectra of the barely resolved binary, source No. 89, were extracted using the 1 pixel radius. Their signal can be contaminated.

Table <ref> provides an inventory of the key line detections in all sources, while Table <ref> lists the measured H$_2$ line fluxes not corrected for extinction. We include only detections above 3$\sigma$.

He I
Ca I
Mg I
Na I
Si I
K I

2	   	E	   	N	   	N	   	A	   	A	   	A	   	N	   	N	   	N

47	   	E	   	N	   	N	   	N	   	N	   	N	   	N	   	N	   	N

87A	   	E	   	N	   	A	   	A	   	N	   	N	   	N	   	A	   	N

87B	   	E	   	N	   	A	   	A	   	A	   	A?	   	N	   	A	   	N

9	   	E	   	E	   	E	   	A	   	A	   	A	   	N	   	N	   	N

13	   	E	   	E	   	A	   	A?	   	N	   	N	   	N	   	A	   	N

15	   	E	   	E	   	A	   	A	   	A	   	A	   	N	   	A	   	N

26	   	E	   	E	   	?	   	A?	   	N	   	A	   	N	   	E	   	N

33	   	E	   	E	   	E	   	A	   	N	   	N	   	N	   	E	   	N

50	   	E	   	E	   	A	   	A	   	A	   	A	   	A?	   	A	   	N

54	   	E	   	E	   	N	   	A	   	?	   	N	   	N	   	A	   	N

60	   	E	   	E	   	E?	   	A	   	N	   	N	   	N	   	E?	   	N

83	   	E	   	E	   	A	   	A?	   	A?	   	N	   	?	   	A	   	N

111	   	E	   	E	   	E?	   	A	   	A	   	A	   	A?	   	E	   	N

112B	   	E?	   	E	   	A	   	?	   	N	   	N	   	N	   	A	   	N

114	   	E	   	E	   	E	   	?	   	N	   	A?	   	N	   	E	   	N

4	   	E	   	N	   	A?	   	A	   	A	   	A	   	N	   	A	   	A

7	   	E	   	N	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

35	   	E	   	N	   	?	   	A	   	A	   	A	   	N	   	A	   	A

72	   	E	   	?	   	A?	   	A	   	A	   	A	   	A?	   	A	   	A

78	   	E	   	N	   	E	   	A	   	A	   	A	   	A	   	E	   	A

100	   	E	   	N	   	A	   	A	   	A	   	A	   	?	   	A	   	A

49	   	E	   	E	   	E	   	A	   	A	   	A	   	A?	   	E	   	A

73	   	E	   	E	   	E	   	A	   	A	   	A	   	E	   	E	   	A?

91B	   	E	   	E	   	N	   	A	   	A?	   	A	   	A	   	?	   	A

109	   	E	   	E	   	A	   	A	   	A	   	A	   	N	   	A	   	A

16	   	N	   	E	   	N	   	A	   	A	   	A	   	A?	   	N	   	N

24	   	N	   	E	   	N	   	A	   	A	   	A	   	N	   	A	   	N

28	   	N	   	E	   	N	   	A	   	A	   	A	   	A	   	A	   	N

31	   	N	   	E	   	?	   	A	   	N	   	N	   	N	   	A?	   	N

36A	   	N	   	E	   	N	   	A	   	A	   	A	   	A	   	N	   	N

41	   	N	   	E	   	N	   	A	   	A	   	N	   	N	   	A	   	N

56	   	N	   	E	   	A	   	A	   	A	   	A	   	N	   	A	   	N

79	   	N	   	E	   	A	   	A	   	N	   	A	   	N	   	A	   	N

81	   	N	   	E	   	A	   	A	   	A	   	A?	   	N	   	A?	   	N

85	   	N	   	E	   	A	   	A	   	A	   	A	   	N	   	A	   	N

93	   	N	   	E	   	E	   	A	   	N	   	A	   	E?	   	E	   	N

102	   	N	   	E	   	N	   	A	   	A	   	A	   	N	   	N	   	N

118	   	N	   	E	   	E	   	A	   	A	   	A	   	A?	   	E	   	N

3	   	N	   	E	   	N	   	A	   	A	   	A	   	A?	   	A	   	A

6	   	N	   	E	   	N	   	A	   	A	   	A	   	A?	   	A	   	A

10	   	N	   	E	   	E	   	A	   	A	   	A	   	N	   	E?	   	A?

12	   	N	   	E	   	E?	   	A	   	A	   	A	   	A?	   	N	   	A

18	   	N	   	E	   	E	   	A	   	A	   	A	   	A	   	E	   	A

20	   	N	   	E	   	A?	   	A	   	A	   	A	   	A	   	N	   	A

22	   	N	   	E	   	E	   	A	   	A	   	A	   	A?	   	E	   	A

23	   	N	   	E	   	E	   	A	   	A	   	A	   	A?	   	E?	   	A

30	   	N	   	E?	   	E?	   	A	   	A	   	A	   	A?	   	N	   	A

32	   	N	   	E	   	N	   	A	   	A	   	A	   	A	   	N	   	A

34	   	N	   	E	   	E	   	A	   	A	   	A	   	A	   	E?	   	A

38A	   	N	   	E	   	A	   	A	   	A	   	A	   	A?	   	E	   	A

38B	   	N	   	E	   	A	   	A	   	A	   	A?	   	A?	   	E	   	A

48	   	N	   	E	   	E?	   	A	   	A	   	A	   	A	   	N	   	A

52A	   	N	   	E	   	A	   	A	   	A	   	A	   	A	   	A	   	A

52B	   	N	   	E	   	A	   	A	   	A	   	A	   	A	   	A	   	A

55A	   	N	   	E	   	E?	   	A	   	A	   	A	   	A	   	E?	   	A

57	   	N	   	E	   	A?	   	A	   	A	   	A	   	N	   	A?	   	A

80	   	N	   	E	   	A	   	A	   	A	   	A	   	A	   	A	   	A

89A	   	N	   	E?	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

97	   	N	   	E	   	A	   	A	   	A	   	A	   	A?	   	A?	   	A

98	   	N	   	E	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

99	   	N	   	E	   	E	   	A	   	?	   	A	   	E?	   	E?	   	A

5	   	N	   	A	   	N	   	A	   	A	   	A	   	A	   	N	   	A

40	   	N	   	A	   	N	   	A	   	A	   	A	   	A	   	N	   	A

43	   	N	   	A	   	N	   	A	   	A	   	A	   	A	   	A?	   	A

53	   	N	   	A	   	N	   	A	   	A	   	N	   	N	   	A	   	N

61	   	N	   	A	   	E?	   	A	   	A	   	A	   	A	   	E	   	A

63	   	N	   	A	   	E?	   	A	   	A	   	A	   	A	   	E	   	A

65	   	N	   	A	   	N	   	A	   	A	   	A	   	A	   	N	   	A

71	   	N	   	A	   	A	   	A	   	A	   	A?	   	?	   	A	   	A

75	   	N	   	A	   	A	   	A?	   	A	   	N	   	A?	   	A	   	N

84	   	N	   	A	   	N	   	A	   	A	   	A	   	A	   	E?	   	A

91A	   	E	   	A	   	A?	   	A	   	A	   	A	   	A	   	E?	   	A

92	   	N	   	A	   	A	   	A	   	A	   	N	   	N	   	A	   	N

94	   	N	   	A	   	A	   	A	   	A	   	A	   	A?	   	A?	   	A

96	   	N	   	A	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

104	   	N	   	A	   	A?	   	A	   	A	   	A	   	A	   	E?	   	A

116	   	N	   	A	   	A	   	A	   	A	   	A	   	A	   	E	   	A?

19	   	E	   	E	   	E	   	N	   	N	   	N	   	N	   	E	   	E

39	   	E?	   	E	   	N	   	A	   	A	   	A	   	N	   	N	   	E

46	   	E	   	E	   	N	   	A	   	A	   	N	   	N	   	N	   	E

64	   	E	   	E	   	E	   	A	   	A?	   	N	   	N	   	E	   	E

77	   	E	   	E	   	E	   	A?	   	A?	   	A?	   	E?	   	E	   	E

112A	   	E?	   	E	   	A	   	?	   	N	   	N	   	N	   	A	   	E

25	   	N	   	N	   	N	   	A	   	A	   	A	   	A	   	A	   	A

29	   	N	   	N	   	A?	   	A	   	A	   	A	   	A	   	A	   	A

42	   	N	   	N	   	E?	   	A	   	A	   	A	   	N	   	E?	   	A

44	   	N	   	N	   	N	   	A	   	A	   	A	   	A?	   	A?	   	A

45	   	N	   	N	   	N	   	A	   	A	   	A	   	N	   	N	   	A

51	   	N	   	N	   	N	   	A	   	A	   	A	   	A	   	N	   	A

55B	   	N	   	N	   	E	   	A	   	A	   	A	   	A?	   	E?	   	A

58	   	N	   	N	   	E?	   	A	   	A	   	A	   	A?	   	N	   	A

59	   	N	   	N	   	?	   	A	   	A	   	A	   	A?	   	N	   	A

66	   	N	   	N	   	A	   	A	   	A	   	A	   	A	   	N	   	A

68	   	N	   	N	   	N	   	A	   	A	   	A	   	A	   	E?	   	A

69	   	N	   	N	   	E?	   	A	   	A?	   	A?	   	N	   	A	   	A?

70	   	N	   	N	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

74	   	N	   	N	   	E?	   	A	   	A	   	A	   	?	   	A	   	A

82A	   	N	   	N	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

82B	   	N	   	N	   	A	   	A	   	A	   	A	   	A?	   	A	   	A

86A	   	N	   	N	   	N	   	A	   	A	   	A	   	A?	   	A	   	A

86B	   	N	   	N	   	?	   	A	   	A?	   	A	   	A?	   	A?	   	A

89B	   	N	   	N	   	?	   	A	   	A	   	A	   	A?	   	A	   	A

95	   	N	   	N	   	E	   	A	   	A	   	A	   	N	   	E	   	A?

101	   	N	   	N	   	E?	   	A	   	N	   	A	   	E?	   	E	   	A

103	   	N	   	N	   	A?	   	A	   	A	   	A	   	A?	   	E	   	A

105	   	N	   	N	   	E	   	A	   	A	   	A	   	A	   	E	   	A

108	   	N	   	N	   	E	   	A	   	A	   	A	   	A	   	E	   	A

110	   	N	   	N	   	E	   	A	   	A	   	A	   	A?	   	E	   	A

113	   	N	   	N	   	A	   	A	   	A	   	A	   	A	   	A	   	A

115	   	N	   	N	   	A	   	A	   	A	   	A	   	A	   	A	   	A

117	   	N	   	N	   	E	   	A	   	A	   	A	   	A	   	E	   	A

1	   	N	   	N	   	?	   	A	   	A	   	A	   	A	   	E?	   	N

8	   	N	   	N	   	N	   	A	   	A?	   	A?	   	N	   	A	   	N

11	   	N	   	N	   	E	   	A	   	A	   	N	   	N	   	A	   	N

14	   	N	   	N	   	E	   	A?	   	A	   	N	   	E?	   	E	   	N

17	   	N	   	N	   	A	   	A	   	A	   	A	   	N	   	A	   	N

21A	   	N	   	N	   	A?	   	?	   	N	   	N	   	N	   	A	   	N

21B	   	N	   	N	   	A?	   	?	   	A	   	N	   	N	   	A	   	N

62	   	N	   	N	   	E	   	A	   	A	   	A	   	N	   	E	   	N

76	   	N	   	N	   	?	   	N	   	A?	   	N	   	?	   	A?	   	N

90	   	N	   	N	   	A	   	A?	   	A?	   	N	   	N	   	A	   	N

106	   	N	   	N	   	A	   	A	   	N	   	A?	   	N	   	A	   	N

E indicates emission line, A - absorption, N - no detection, ? - possible detection.

lcccccccBlccccccc 86.8

H$_2$ Emission Line Fluxes.

No.     1-0 S(0)     1-0 S(1)
1-0 S(2)     1-0 S(3)
2-1 S(1)     2-1 S(2)
2-1 S(3)
1-0 Q(1)     1-0 Q(2)
1-0 Q(3)     1-0 Q(4)
1-0 Q(5)     1-0 Q(6)

2-8 10-15
7c($\rm 10^{-16}~erg~s^{-1}cm^{-2}$)          6c($\rm 10^{-16}~erg~s^{-1}cm^{-2}$)

2     5.35 $\pm$ 0.33     25.61 $\pm$ 0.58     10.18 $\pm$ 0.38     36.90 $\pm$ 3.39             2.24 $\pm$ 0.87     2     25.77 $\pm$ 1.54     5.97 $\pm$ 1.86     22.60 $\pm$ 1.10         16.72 $\pm$ 1.14     5.40 $\pm$ 1.37

4*     1.55 $\pm$ 0.16     3.36 $\pm$ 0.20     0.85 $\pm$ 0.11     6.27 $\pm$ 1.19     0.97 $\pm$ 0.26             4*     5.38 $\pm$ 0.84         4.99 $\pm$ 0.65         2.98 $\pm$ 0.56

7*         0.87 $\pm$ 0.15                         7*                     2.57 $\pm$ 0.76

9         1.62 $\pm$ 0.12         1.83 $\pm$ 0.51                 9

13         2.33 $\pm$ 0.56         3.26 $\pm$ 0.93                 13

15         4.42 $\pm$ 0.27         7.84 $\pm$ 2.26     1.91 $\pm$ 0.31             15     5.70 $\pm$ 1.22         6.59 $\pm$ 1.45         4.21 $\pm$ 0.92

19     1.52 $\pm$ 0.25     5.30 $\pm$ 0.22     2.62 $\pm$ 0.24     5.14 $\pm$ 0.84     1.21 $\pm$ 0.31             19     6.21 $\pm$ 0.83         6.71 $\pm$ 0.90         3.62 $\pm$ 1.02     2.24 $\pm$ 0.67

26         1.51 $\pm$ 0.20     0.86 $\pm$ 0.20                     26

33     0.94 $\pm$ 0.30     2.17 $\pm$ 0.16     0.83 $\pm$ 0.15     5.43 $\pm$ 1.39                 33     3.53 $\pm$ 0.75         3.30 $\pm$ 1.00

35         3.24 $\pm$ 0.31                         35

39         12.53 $\pm$ 1.57                         39             22.78 $\pm$ 2.78

46         12.92 $\pm$ 0.60     6.16 $\pm$ 0.81                     46

47     17.48 $\pm$ 0.52     69.91 $\pm$ 1.05     22.15 $\pm$ 0.51     76.05 $\pm$ 2.38     7.43 $\pm$ 0.40     3.09 $\pm$ 0.36     6.19 $\pm$ 0.34     47     85.63 $\pm$ 2.80     24.76 $\pm$ 0.95     76.93 $\pm$ 2.74     20.75 $\pm$ 3.33     26.15 $\pm$ 1.24     10.96 $\pm$ 2.36

49         4.65 $\pm$ 0.34     3.14 $\pm$ 0.47                 3.30 $\pm$ 0.59     49     6.04 $\pm$ 0.53         4.87 $\pm$ 0.53

50         1.47 $\pm$ 0.18     0.49 $\pm$ 0.17     4.37 $\pm$ 1.22                 50             2.09 $\pm$ 0.56         4.97 $\pm$ 1.02

54*     23.05 $\pm$ 6.16     75.32 $\pm$ 14.15         123.66 $\pm$ 27.71                 54*

60     14.58 $\pm$ 0.69     60.13 $\pm$ 0.47     23.57 $\pm$ 0.71     99.90 $\pm$ 13.36     7.30 $\pm$ 0.79     2.74 $\pm$ 0.36     5.84 $\pm$ 1.29     60     65.22 $\pm$ 3.30     12.80 $\pm$ 1.80     58.16 $\pm$ 2.53     39.66 $\pm$ 9.77     24.78 $\pm$ 5.70     13.01 $\pm$ 4.73

64     3.11 $\pm$ 0.42     9.67 $\pm$ 0.45     4.86 $\pm$ 0.23             0.79 $\pm$ 0.25     2.45 $\pm$ 0.77     64     10.93 $\pm$ 1.08     4.87 $\pm$ 1.14     11.93 $\pm$ 0.62             3.51 $\pm$ 1.05

73         1.64 $\pm$ 0.32     2.45 $\pm$ 0.26             0.69 $\pm$ 0.13         73     2.78 $\pm$ 0.58         2.49 $\pm$ 0.33         2.37 $\pm$ 0.53

77         1.37 $\pm$ 0.14     2.43 $\pm$ 0.26         0.40 $\pm$ 0.14             77     1.64 $\pm$ 0.46

83         2.88 $\pm$ 0.26         3.25 $\pm$ 1.01                 83

87A     0.49 $\pm$ 0.17     2.11 $\pm$ 0.19                         87A     1.52 $\pm$ 0.52         1.17 $\pm$ 0.33

87B         0.66 $\pm$ 0.08                         87B     0.40 $\pm$ 0.17

91A*         11.73 $\pm$ 1.70         17.49 $\pm$ 4.68         5.43 $\pm$ 1.05         91A*     20.27 $\pm$ 6.68         17.11 $\pm$ 4.79

91B*     1.21 $\pm$ 0.32     4.84 $\pm$ 0.26     1.93 $\pm$ 0.32     4.14 $\pm$ 0.56             0.79 $\pm$ 0.23     91B*     5.87 $\pm$ 0.50         4.77 $\pm$ 0.25

100*     2.56 $\pm$ 0.28     10.07 $\pm$ 0.25     3.77 $\pm$ 0.54     6.33 $\pm$ 0.92     1.40 $\pm$ 0.25             100*     8.61 $\pm$ 0.48     2.63 $\pm$ 0.67     6.22 $\pm$ 0.38

109     2.12 $\pm$ 0.18     9.18 $\pm$ 0.46     2.61 $\pm$ 0.27     7.89 $\pm$ 1.44     0.80 $\pm$ 0.14             109     9.26 $\pm$ 0.42     2.91 $\pm$ 0.56     6.85 $\pm$ 0.66

111         3.93 $\pm$ 0.90     2.14 $\pm$ 0.49                     111

112A     0.27 $\pm$ 0.03     0.79 $\pm$ 0.16                         112A             1.73 $\pm$ 0.34

112B         0.58 $\pm$ 0.15     0.25 $\pm$ 0.09                     112B

114     2.64 $\pm$ 0.31     10.46 $\pm$ 0.65     4.93 $\pm$ 0.19     6.72 $\pm$ 1.08     0.83 $\pm$ 0.25     1.69 $\pm$ 0.40     2.97 $\pm$ 0.69     114     12.57 $\pm$ 1.00         9.87 $\pm$ 0.79

H$_2$ line fluxes have not been corrected for extinction. Reported are only fluxes with detection significance above 3$\sigma$. Targets with asterisk do not have estimated extinction from SED models.

$K$-band spectra of YSO candidates in the CMa-$\ell$224 star forming region, which show the H Br$\gamma$ line in emission. All spectra are continuum subtracted and normalised to the peak flux density in the range from 1.97 to 2.47 $\mu$m. Red horizontal lines show spectral ranges most affected by telluric lines; the range at 2.0 $\mu$m has been removed for clarity.

Spectra with Br$\gamma$ and H$_2$ emission lines.

Spectra with only H$_2$ emission lines.

Spectra with Br$\gamma$ and H$_2$  emission lines, and CO bandhead in absorption.

Spectra with Br$\gamma$ emission line and CO bandhead in absorption.

Spectra with H$_2$ emission lines and CO bandhead in absorption.

Spectra with Br$\gamma$ absorption line.

Spectra with CO bandhead in emission.

Spectra with CO bandhead in absorption.

Spectra with no emission lines and no CO or Br$\gamma$ absorption lines.

Spectra from three pairs of double sources.

Spectra from double sources - continuation.

§ LINE EMISSION MAPS

Figures <ref>–<ref> show the emission maps of the most commonly detected H$_2$ lines in the $K$-band range. Figure <ref> shows the spatial extent of several transitions of
H$_2$ toward source No. 47, which has the most molecular-rich spectrum at $K$-band. Source No. 47 has the K-band magnitude of 15.2 based on our KMOS observations, thus it is one of the faintest sources in the sample. [Sewiło et al., 2019] classified it as a Class I YSO. The young evolutionary stage of the source is confirmed by strong bipolar outflows. This source is an interesting case for detailed follow-up studies.

Figure <ref> shows the hydrogen Br$\gamma$ 2.17 $\mu$m emission maps. Figures <ref>, <ref>, <ref>, and <ref> show the emission maps for several CO bandhead transitions – CO 2-0, 3-1, 4-2, and CO 5-3, respectively.

Integrated maps of the H$_2$ 1-0 S(1) line at 2.1218 $\mu$m for a subsample of sources where the extended H$_2$ emission is detected (see Table <ref>). Colors correspond to the line fluxes at each pixel (detections above 3$\sigma$) and white contours correspond to the continuum emission in $K$-band (see also Appendix A). Source names are provided in the upper right corners (No. in Table <ref>). Note that source No. 88 (G224.4514$-$00.7054) has a non-detection of the continuum emission at 2.1218 $\mu$m.

Similar to Figure <ref>, but for the H$_2$ 1-0 S(0) line at 2.2235 $\mu$m. Only emission above 2$\sigma$ is shown.

Similar to Figure <ref>, but for the H$_2$ 1-0 Q(1) line at 2.4066 $\mu$m.   Only emission above 2$\sigma$ is shown.

Similar to Figure <ref>, but for the H$_2$ 1-0 Q(3) line at 2.4237 $\mu$m.   Only emission above 2$\sigma$ is shown.

Similar to Figure <ref>, but for the H$_2$ 2-1 S(1) line at 2.2477 $\mu$m.   Only emission above 2$\sigma$ is shown.

Maps in several H$_2$ lines in source No. 47, which shows detections of the highest-excited H$_2$ lines in the entire sample (see Table <ref>).

Integrated maps for the H Br$\gamma$ line at 2.1661 $\mu$m. Colors correspond to the line fluxes at each pixel (detections above 2$\sigma$) and white contours show the continuum emission in $K$-band (see also Appendix <ref>).

Integrated maps of the CO $\nu$ = 2-0 line at 2.29 $\mu$m. Colors correspond to the line fluxes at each pixel (detections above 2$\sigma$) and white contours show the continuum emission in K-band (see also Appendix <ref>).

Similar to Figure <ref>, but the maps show the CO $\nu$ = 3-1 line at 2.32 $\mu$m.

Similar to Figure <ref>, but the maps show the CO $\nu$ = 4-2 line at 2.35 $\mu$m.

Similar to Figure <ref>, but the maps show the CO $\nu$ = 5-3 line at 2.38 $\mu$m.

§ EXTINCTION WITHIN STELLAR ENVELOPES

The extinction values from SED modelling incorporate only the interstellar extinction to the source. Some of our stars were identified to be embedded in the envelope. Thus, to complete the estimations of extinction,
we investigate the level of extinction due to the envelope.
We consider only the targets, for which best-fit SED models included the envelope component [Sewiło et al., 2019]. For simplicity, we assume that the spherical envelope can be well described by the power-law density profile <cit.>:

\begin{equation}
\rho(r) = \rho_0 \cdot (r/R_0)^\gamma
\end{equation}

where $\rho_0$ is a volume mass density at the radius $R_0$ and $\gamma$ represents the slope of the relation. We do not include outflow cavities, which were considered in the original models.

To calculate column densities, we integrate the density profile along the line of sight, between the inner and outer radii of the envelope. We assume the mean molecular weight of 2.8 and the gas-to-dust ratio of 100. For calculations, we adopt parameter values from SED modelling of [Sewiło et al., 2019] and set the outer radius to be at 5000 au (the maximum value considered in models of Robitaille, 2017). In cases where the inner radius was not available, we assume a value of 50 au. For target No. 47 we also assume the $R_0=100$ au, which is within the parameter range adopted in the SED modelling [Robitaille, 2017]. Some of our sources were originally fitted by a broken power-law density profile <cit.> and therefore they lack the value of $\gamma$. We assume a slope of -1/2 for $r<R_0$ and -3/2 for $r>R_0$, following the description in Sec. 3.2.3 of [Robitaille, 2017].

We estimate the visual extinction due to the envelope of YSOs based on derived column densities. We follow the formula from [Predehl & Schmitt, 1995]:

\begin{equation}
A_{\rm V} = 1.086 \cdot N \cdot \pi a_g^2 \cdot Q_{\rm ext}
\end{equation}

where $N$ is a column density in 10$^{21}$ cm$^{-2}$, $a_g$ is a radius of a spherical grain in 0.1 $\mu$m, and $Q_{\rm ext}$ is an extinction efficiency. We adopted $a_g$ = 0.1 $\mu$m and $Q_{\rm ext}$ = 0.5. The resulting estimates are shown in Fig. <ref> as a function of a total $A_{\rm V}$ obtained from the SED models. $A_{\rm{V,env}}$  is higher for most Class I than for Class II sources, indicating that envelope is being dissipated with time, in line with evolutionary predictions. The only outlier, Class II star No. 65, is characterised by high $\rho_0 = 1.8\times10^{-18}$ g cm$^{-3}$, which leads to the very high value of the envelope extinction. This source is very bright, has relatively low $A_{\rm V}$, and has no emission lines, which are signposts of ongoing accretion or ejection. Clear detection of CO bandhead in absorption suggests that this target might be a foreground contamination, possibly a red giant.

Interstellar and Envelope Extinction Values from SED Modeling.

No.     $A_{\rm{V}}$ (Q2/S0)     $A_{\rm{V}}$ (Q3/S1)     $A_{\rm{V}}$ (Q4/S2)     $A_{\rm{V}}$ (Q5/S3)     $A_{\rm{V,interst}}$     $A_{\rm{V,env}}$

1                     8.04

2     0.80 $\pm$ 0.26     9.75 $\pm$ 0.53         -1.51 $\pm$ 0.17     9.0     12.99

3                     11.9

5                     34.9     3.98

6                     23.4

8                     25.3

9     81.50 $\pm$ 49.80                 30.8

10                     27.6

12                     31.0

13                     15.4

15     29.52 $\pm$ 18.11     32.04 $\pm$ 7.30         2.52 $\pm$ 0.91     40.0

16                     21.3

19     -28.56 $\pm$ 27.09     25.06 $\pm$ 3.51         8.92 $\pm$ 2.90     33.9

20                     12.2

24                     27.0

26                     25.9

28                     9.92

30                     2.84

32                     6.84

33         32.84 $\pm$ 10.22             21.7

35                     20.0     2.06

39         40.42 $\pm$ 7.08     49.49 $\pm$ 26.88         10.2

40                     28.6

42                     26.2

43                     2.29

45                     19.3

46                     14.0

47     17.83 $\pm$ 0.86     19.12 $\pm$ 0.74     15.79 $\pm$ 2.56     -8.00 $\pm$ 0.46     33.3     12.65

48                     3.03

49         17.03 $\pm$ 2.23     14.59 $\pm$ 14.10         32.0

50         30.12 $\pm$ 8.94         20.18 $\pm$ 7.00     15.7

53                     22.4

55A                     3.37

55B                     3.37

56                     22.0

58                     5.03

59                     2.9

60     -16.33 $\pm$ 2.42     13.64 $\pm$ 0.60     33.50 $\pm$ 8.31     -15.68 $\pm$ 4.17     18.2     22.21

61                     8.51

62                     17.3

63                     22.6

64     24.99 $\pm$ 6.75     23.97 $\pm$ 1.67     38.78 $\pm$ 27.31         7.15     28.74

65                     3.84     109.96

68                     7.19

70                     10.6

71                     27.8

73         32.84 $\pm$ 7.77         30.29 $\pm$ 26.38     26.2

75                     23.1

76                     24.2

77                     40.0     0.00

79                     15.1

81                     25.7

82A                     11.8     0.01

82B                     11.8     0.01

83         -9.06 $\pm$ 7.47             12.8     41.15

84                     17.3     3.13

86A                     6.14

86B                     6.14

87A     94.83 $\pm$ 70.51     -9.75 $\pm$ 2.87             30.3

87B     77.67 $\pm$ 61.67             49.16 $\pm$ 43.48     30.3

88                     0.0     0.23

92                     9.37

93                     28.5

94                     8.96

95                     4.49

96                     19.1     0.04

97                     13.9

98                     7.47

102                     15.4

104                     7.78

105                     28.1

108                     7.74

109     15.66 $\pm$ 3.31     2.66 $\pm$ 0.29             15.0     29.33

110                     16.9

111     50.66 $\pm$ 45.27     16.41 $\pm$ 6.99             17.0

112A         48.09 $\pm$ 13.35             25.1

112B                     25.1

113                     21.7

114     -3.26 $\pm$ 1.76     12.61 $\pm$ 1.27             39.4

115                     9.23     15.27

116                     2.43

117                     5.11

118                     12.7

All values are given in magnitudes. $A_{\rm V,interst}$
is the value adopted from [Sewiło et al., 2019] based on the SED modeling; we assume an uncertainty of 10 mag of their measurements (see text for details). The $A_{\rm V,env}$ is the extinction within the envelope estimated based on the results of the SED modeling.

Comparison between the total visual extinction $A_{\rm{V,total}}$ [Sewiło et al., 2019] and the extinction within the envelope $A_{\rm{V,env}}$. The latter was estimated assuming spherical envelope and a power-law density profile. See the text for details.