# Investigating Protostellar Accretion across the mass spectrum with the JWST:
discovery of a collimated jet from the low luminosity protostar IRAS
16253-2429 in a quiescent accretion phase

Mayank Narang Academia Sinica Institute of Astronomy & Astrophysics, 11F of
Astro-Math Bldg., No.1, Sec. 4, Roosevelt Rd., Taipei 10617, Taiwan, R.O.C.
Department of Astronomy and Astrophysics Tata Institute of Fundamental
Research
Homi Bhabha Road, Colaba, Mumbai 400005, India  P. Manoj Department of
Astronomy and Astrophysics Tata Institute of Fundamental Research
Homi Bhabha Road, Colaba, Mumbai 400005, India  Himanshu Tyagi Department of
Astronomy and Astrophysics Tata Institute of Fundamental Research
Homi Bhabha Road, Colaba, Mumbai 400005, India  Dan M. Watson University of
Rochester, Rochester, NY, US  S. Thomas Megeath University of Toledo, Toledo,
OH, US  Samuel Federman University of Toledo, Toledo, OH, US  Adam E.
Rubinstein University of Rochester, Rochester, NY, US  Robert Gutermuth
University of Massachusetts Amherst, Amherst, MA, US  Alessio Caratti o
Garatti INAF-Osservatorio Astronomico di Capodimonte, IT  Henrik Beuther Max
Planck Institute for Astronomy, Heidelberg, Baden Wuerttemberg, DE  Tyler L.
Bourke SKA Observatory, Jodrell Bank, Lower Withington, Macclesfield SK11 9FT,
UK  Ewine F. Van Dishoeck Leiden Observatory, Universiteit Leiden, Leiden,
Zuid-Holland, NL  Max-Planck Institut für Extraterrestrische Physik, Garching
bei München, DE  Neal J. Evans II Department of Astronomy, The University of
Texas at Austin, 2515 Speedway, Stop C1400, Austin, Texas 78712-1205, USA
Guillem Anglada Instituto de Astrofísica de Andalucía (IAA-CSIC), Granada,
Andalucía, ES  Mayra Osorio Instituto de Astrofísica de Andalucía (IAA-CSIC),
Granada, Andalucía, ES  Thomas Stanke Max-Planck Institut für
Extraterrestrische Physik, Garching bei München, DE  James Muzerolle Space
Telescope Science Institute, Baltimore, MD, US  Leslie W. Looney Department of
Astronomy, University of Illinois, 1002 West Green St, Urbana, IL 61801, USA
National Radio Astronomy Observatory, 520 Edgemont Rd., Charlottesville, VA
22903 USA  Yao-Lun Yang RIKEN Cluster for Pioneering Research, Wako-shi,
Saitama, 351-0106, Japan  Pamela Klaassen United Kingdom Astronomy Technology
Centre, Edinburgh, GB  Nicole Karnath Space Science Institute, Boulder, CO, US
Center for Astrophysics Harvard & Smithsonian, Cambridge, MA, US  Prabhani
Atnagulov Ritter Astrophysical Research Center, Dept. of Physics and
Astronomy, University of Toledo, Toledo, OH, US  Nashanty Brunken Leiden
Observatory, Universiteit Leiden, Leiden, Zuid-Holland, NL  William J. Fischer
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218,
US  Elise Furlan Caltech/IPAC, Pasadena, CA, US  Joel Green Space Telescope
Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, US  Nolan Habel
Jet Propulsion Laboratory, Pasadena, CA, US  Lee Hartmann University of
Michigan, Ann Arbor, MI, US  Hendrik Linz Max Planck Institute for Astronomy,
Heidelberg, Baden Wuerttemberg, DE  Friedrich-Schiller-Universität, Jena,
Thüringen, DE  Pooneh Nazari Leiden Observatory, Universiteit Leiden, Leiden,
Zuid-Holland, NL  Riwaj Pokhrel Ritter Astrophysical Research Center, Dept. of
Physics and Astronomy, University of Toledo, Toledo, OH, US  Rohan
Rahatgaonkar Gemini South Observatory, La Serena, CL  Will R. M. Rocha
Laboratory for Astrophysics, Leiden Observatory, Universiteit Leiden, Leiden,
Zuid-Holland, NL  Patrick Sheehan National Radio Astronomy Observatory, 520
Edgemont Rd., Charlottesville, VA 22903 USA  Katerina Slavicinska Leiden
Observatory, Universiteit Leiden, Leiden, Zuid-Holland, NL  Amelia M. Stutz
Departamento de Astronomía, Universidad de Concepción,Casilla 160-C,
Concepción, Chile  John J. Tobin National Radio Astronomy Observatory, 520
Edgemont Rd., Charlottesville, VA 22903 USA  Lukasz Tychoniec European
Southern Observatory, Garching bei München, DE  Scott Wolk Center for
Astrophysics Harvard & Smithsonian, Cambridge, MA, US

###### Abstract

Investigating Protostellar Accretion (IPA) is a JWST Cycle 1 GO program that
used NIRSpec IFU and MIRI MRS to obtain 2.9–28 $\mu$m spectral cubes of young,
deeply embedded, protostars with luminosities of 0.2 to 10,000 L⊙ and central
masses of 0.15 to 12 M⊙. In this Letter, we report the discovery of a highly
collimated atomic jet from the Class 0 protostar IRAS 16253-2429, the lowest
luminosity (mass) source in the IPA program. The collimated jet is clearly
detected in multiple [Fe II] lines, [Ne II], [Ni II], and H I lines, but not
in molecular emission. From our analysis, we find that the atomic jet has a
velocity of about 169 $\pm$ 15 km s-1, after correcting for inclination. The
width of the jet increases with distance from the central protostar from 23 to
60 au, corresponding to an opening angle of 2.6 $\pm$ 0.5°. By comparing the
measured flux ratios of various fine structure lines to those predicted by
simple shock models, we derive a shock velocity of 54 km s-1 and a preshock
density of 2.0$\times 10^{3}$ cm-3 at the base of the jet. From these
quantities, and assuming a cylindrical cross-section for the jet, we derive an
upper limit for the mass loss rate from the protostar of 1.1$\times
10^{-10}~{}M_{\odot}$ yr -1. The low mass loss rate is consistent with
simultaneous measurements of low mass accretion rate ($\sim 1.3\times
10^{-9}~{}M_{\odot}$ yr -1) for IRAS 16253-2429 from JWST observations (Watson
et al. in prep), indicating that the protostar is in a quiescent accretion
phase. Our results demonstrate that very low-mass protostars can drive highly
collimated, atomic jets, even during the quiescent phase.

##  1 Introduction

Jets and outflows from protostars and pre-main sequence stars play a crucial
role in the process of star formation. These dynamic phenomena are
instrumental in facilitating the removal of angular momentum from the
surrounding disk material, thereby enabling efficient accretion of mass onto
the central object (e.g., Shang et al., 2007; Pudritz et al., 2007; Hartmann
et al., 2016; Frank et al., 2014; Bally, 2016; Ray & Ferreira, 2021). These
accretion-driven outflows/jets also play vital roles in shaping the Initial
Mass Function (IMF) (e.g., Bate, 2009; Hennebelle et al., 2020; Appel et al.,
2023) and can inject energy and momentum into the surroundings, and disperse a
significant fraction of the protostellar envelope, which limits star formation
efficiency (e.g., Bate, 2009; Fall et al., 2010; Hennebelle et al., 2020;
Frank et al., 2014).

Figure 1: (a) The Spitzer IRAC 3 color (IRAC 1, 2, and 4 are R G B
respectively) image of the IRAS 16253-2429 field. The field of view is
2′$\times$ 2′. The black square in the center represents the JWST FOV. (b) A
zoomed-in view of the NIRSpec and MIRI channel-1 short field with the red
representing continuum emission at 4.17 $\mu$m, the green showing H2 0-0 S(11)
at 4.18 $\mu$m and the blue depicting [Fe II] at 5.34 $\mu$m. A local
continuum has been subtracted to make the line maps.

Despite their unquestionable significance, the precise origins of protostellar
jets and outflows remain uncertain and continue to be a subject of debate.
Among the various proposed mechanisms, the most widely accepted one is the
magnetocentrifugal acceleration (Blandford & Payne, 1982; Pudritz & Norman,
1983; Pelletier & Pudritz, 1992; Wardle & Koenigl, 1993; Najita & Shu, 1994;
Shu et al., 1994; Matt & Pudritz, 2005, 2008; Frank et al., 2014). According
to this model, material situated near the surface of the inner parts of the
disk is ejected centrifugally along the rotating magnetic field lines. As the
outflow progresses beyond the so-called “Alfvén surface,” where the kinetic
energy matches the magnetic energy, the dynamics begin to dominate over the
magnetic field, resulting in a predominantly toroidal magnetic configuration.
The tension force generated by these magnetic loops focuses the outflowing
material into a narrow jet along the rotational axis of the circumstellar disk
that originally launched the (MHD) wind (Blandford & Payne, 1982; Spruit,
1996; Ray & Ferreira, 2021). As a result, highly collimated bipolar jets are
formed.

The three classes of models that are most successful in explaining the
launching of outflows from YSOs are X-wind, the disk-wind, and accretion-
powered stellar winds. In the X-wind model (e.g. Shu et al., 1988, 2000), the
jet-launching magnetic field is concentrated near the inner edge of the disk,
known as the ”X-point,” where the stellar magnetosphere truncates the disk.
The disk-wind model (e.g. Pelletier & Pudritz, 1992; Pudritz et al., 2007)
posits that the outflow-launching field lines (threaded through the disk)
extend over a broader range of larger disk radii (also see Lee et al. 2022).
In the accretion-powered stellar wind scenario (e.g., Matt & Pudritz, 2005,
2008), a fraction of the energy released by mass accreting onto the protostar
powers the stellar wind (also see Watson et al. 2016). Though all these models
predict that the mass ejection rate is proportional to the mass accretion
rate, the constant of proportionality depends upon the underlying model
(Watson et al., 2016). Therefore to better understand the launching of the
jets/outflow from protostars, it is critically important to study the mass
ejection and accretion rates from them.

Figure 2: The morphology of [Fe II] lines detected towards IRAS 16253-2429.
All these lines trace the collimated jet launched from the protostar. All
images are cropped to the same spatial scale. The white solid circle is the
MIRI MRS 14 $\mu$m continuum position. The scale bar corresponding to 100 au
is shown in the bottom left corner.

Infrared emission lines of [Fe II] and H2 are well-known tracers of shocked
gas and have been used to trace outflows and jets from protostars (e.g.
Caratti o Garatti et al., 2006; Agra-Amboage et al., 2014; Reipurth et al.,
2019; Watson et al., 2016; Bally, 2016; Yang et al., 2022; Beuther et al.,
2023; Harsono et al., 2023; Rubinstein et al., 2023; Ray, T. P. et al., 2023).
These lines serve as excellent diagnostics for deriving the physical
conditions of the collimated jets as well as the shocked molecular gas in the
outflow (Ray, T. P. et al., 2023). Therefore by harnessing the spatial and
spectral resolution as well as the sensitivity of JWST, we can gain a deeper
understanding of the launching of and feedback from the protostellar
outflows/jets.

Table 1: The log of our JWST observations. Observation | RA | Dec | Hours | Start UT | End UT
---|---|---|---|---|---
NIRSpec IFU IRAS 16253-2429 | 16 28 21.6200 | -24 36 24.16 | 3.43 | Jul 22, 2022 15:09:15 | Jul 22, 2022 17:59:43
MIRI MRS IRAS 16253-2429 | 16 28 21.6200 | -24 36 24.16 | 7.43 | Jul 23, 2022 04:04:41 | Jul 23, 2022 09:58:57
MIRI MRS IRAS 16253-2429 background | 16 28 27.7560 | -24 37 46.95 | 1.86 | Jul 23, 2022 10:02:33 | Jul 23, 2022 11:35:19

Investigating Protostellar Accretion (IPA) across the mass spectrum is a JWST
Cycle 1 medium GO program (Program ID 1802, PI Tom Megeath; Megeath et al.
2021, Federman et al. 2023b Rubinstein et al., in prep, Watson et al., prep).
The IPA program aims to study the processes that drive protostellar accretion,
and accretion-driven jets/outflows as a function of protostellar mass. The
program also aims to characterize the chemical and spatial variation of ices
in the surrounding protostellar envelope (Brunken et al., in prep, Tyagi et
al., in prep, Nazari et al., in prep and Slavicinska et al., in prep). To
achieve this, the IPA program utilizes the NIRSpec+MIRI Integral Field Units
(IFUs) to obtain spectral cubes spanning the wavelength range of 2.9-28
$\mu$m. The IPA protostellar sample consists of five sources that range in
bolometric luminosities from 0.2 to 10,000 $L_{\odot}$ and in central masses
from 0.15 to 12 $M_{\odot}$. These protostars are in the Class 0 phase, during
which most of the stellar mass is accreted (e.g., Fischer et al., 2017;
Federman et al., 2023a; Narang et al., 2023).

The protostar IRAS 16253-2429 is the lowest mass and lowest luminosity source
in the IPA sample. It is located in a relatively isolated region of the
$\rho$-Ophiuchus molecular cloud (Khanzadyan et al., 2004; Stanke et al.,
2006; Barsony et al., 2010). With a bolometric luminosity of only 0.2
$L_{\odot}$ and a bolometric temperature of 42 K (Pokhrel et al., in
preparation; also see Pokhrel et al. 2023), coupled with an internal
luminosity below 0.15 $L_{\odot}$ (Pokhrel et al., in preparation), IRAS
16253-2429 falls under the category of a very low luminosity object (VeLLO)
(Dunham et al., 2008). Being only at a distance of 140 pc (average of
$\rho$-Ophiuchi molecular cloud members Zucker et al. 2020) and relatively
isolated, IRAS 16253-2429 is an ideal target to study the morphology and
kinematics of jets from very low mass ($<0.2$ $M_{\odot}$) protostars.

Extensive studies have been conducted on this protostar, particularly at the
(sub)millimeter wavelengths (Stanke et al., 2006; Tobin et al., 2011; Yen et
al., 2017; Hsieh et al., 2019). These observations have revealed the presence
of a central protostar associated with a bipolar outflow identified in CO, as
well as a Keplerian disk. However, no collimated molecular jet in CO or SiO
has been detected from the system (Aso et al., 2023). This is consistent with
the low detection rate of molecular jets from low luminosity protostars (e.g.,
Podio et al., 2021). Recent observations with ALMA have estimated the
dynamical mass of the central source to be between 0.12-0.17 $M_{\odot}$ (Aso
et al., 2023), indicating that IRAS 16253-2429 has already accreted a mass
above the hydrogen-burning limit. The protostar also has a tiny disk with a
radius of 16 $\pm$ 3 au (Aso et al., 2023) with a mass of $\sim$ 2 $\times
10^{-3}\,M_{\odot}$. The envelope mass of IRAS 16253-2429 also has been
estimated from mm observations are in the range of 0.2 to 1 $M_{\odot}$
(Stanke et al., 2006; Enoch et al., 2008; Tobin et al., 2011). The disk-to-
envelope mass ratio for IRAS 16253-2429 therefore is $\ll$ 0.01 and is
consistent with the values reported for Class 0 protostars by Federman et al.
(2023a).

Figure 1 shows the Spitzer Infrared Array Camera (IRAC) images of the IRAS
16253-2429 field, which reveal a bipolar hourglass structure in the Northeast
to Southwest direction. This distinctive shape is traced by the outflow
cavities in H2 and the continuum. The bipolar outflow from the source extends
up to 2′in size, corresponding to a physical scale of 16800 AU or 0.08 pc. In
addition to the imaging, the Spitzer IRS instrument also carried out spectral
scan mapping of the protostar (Barsony et al., 2010). The Spitzer IRS spectra
detected six pure rotational H2 (0-0 S(2) to 0-0 S(7)) lines along the
protostellar outflow. However, no fine structure lines were detected from the
protostar with Spitzer/IRS.

##  2 Observations and Data reduction

Our primary objective with the JWST observations was to perform spectral
imaging by combining the Integral Field Unit (IFU) observations from the Mid-
Infrared Instrument (MIRI) Medium Resolution Spectrograph (MRS) (Rieke et al.,
2015; Wright et al., 2015) and the Near Infrared Spectrograph (NIRSpec)
(Jakobsen et al., 2022; Böker et al., 2022) covering from 2.87 to 28 $\mu$m.
The NIRSpec IFU and MIRI MRS observations were conducted near-simultaneously.
The details of the observation log are provided in Table 1.

For the NIRSpec IFU observations, we utilized the F290L/G395M filter/grating
combination. To ensure comprehensive coverage, we employed a $2\times 2$
mosaic pattern with a 10% overlap and a 4-point dither mode. The NIRSpec field
of view (FOV) is $\sim$ 6″$\times$ 6″. The total exposure time, which
accounted for overheads, amounted to 3.43 hours (for further detail, see
Federman et al. 2023b).

Figure 3: The morphology Br-$\alpha$ line and the other atomic fine structure
lines and detected towards IRAS 16253-2429. All images are cropped to the same
spatial scale. The white solid circle is the MIRI MRS 14 $\mu$m continuum
position. The scale bar corresponding to 100 au is shown in the bottom right
corner.

IRAS 16253-2429 was observed using all four channels of the MIRI MRS. Similar
to the NIRSpec observations, we employed a $2\times 2$ mosaic pattern with a
10% overlap and a 4-point dither for the MIRI observations. The total exposure
time, including overheads, for the science target amounted to 7.43 hours.
Additionally, a dedicated background observation with a similar setup and an
exposure time of 1.86 hours (including overheads) was obtained. The dedicated
background field was positioned approximately 118″away from the target. The
field of view (FOV) for MIRI varies from channel to channel with channel-1
short covering a FOV of $\sim$ 6″$\times$ 6″while for channel-4 long the FOV
is $\sim$ 15″$\times$ 15″. The orientation and FOV of the MIRI MRS is also
different from NIRSpec IFU.

We reduced the NIRSpec IFU using the JWST pipeline version 1.9.5 and the JWST
Calibration References Data System context version jwst_1069.pmap. Our
analysis of the NIRSpec data showed the presence of several hot pixels that
were being missed by the JWST outlier detection step. These hot pixels were a
function of both wavelengths as well as time. Hence a universal dark mask
could not be used. To remedy this, we used a custom outlier detection
algorithm for the NIRSpec observations to remove the hot pixels (see Federman
et al. 2023b). In addition, MIRI parallel imaging observations were obtained
at both MIRI pointings.

Figure 4: (a) The horizontal slices in green marked on top of the jet traced
in [Fe II] at 5.34 $\mu$m. The white circle represents the protostellar
position. (b) The width (FWHM) of the jet as a function of distance from the
central protostar. The negative distance corresponds to the northern side, and
the positive distance corresponds to the southern side. (c) The deconvolved
width (FWHM) of the jet as a function of distance from the central protostar.
The fit to the jet width as a function of distance is shown as the red solid
line. The green triangles is sub figure (b) and (c) show the upper limits to
the width at those points.  Figure 5: The moment 1 map for the [Fe II] line at
5.34 $\mu$m. The white solid circle is the MIRI MRS 14 $\mu$m continuum
position. The scale bar corresponding to 100 au is shown in the bottom left
corner.

To reduce the MIRI data we used the JWST pipeline version 1.11.3 and the JWST
Calibration References Data System context version jwst_1100.pmap. We used the
standard Stage-1 JWST pipeline Detector1Pipeline to reduce the MIRI MRS data
starting from uncal data. In Stage-2 (Spec2Pipeline), the dedicated background
was subtracted, and the residual fringe correction was also performed. The
Stage-3 Spec3Pipeline step was then run with the CubeBuildStep set to ’band’
mode such that each channel and each band are reduced as a separate fits file.
However, on examining the dedicated background IFU image, we detected emission
from the H2 S (1) (17.035 $\mu$m) and H2 S(2) (12.279 $\mu$m) molecular lines.
Therefore we also reduced the MIRI MRS data without subtracting the dedicated
background. We will use these non-background subtracted data cubes for
measuring the extinction in Section 3.3.

During our analysis of the data obtained from the MIRI MRS and NIRSpec IFU
instruments, we found a positional offset between the two instruments. To
accurately determine and quantify this offset between the two instruments, we
focused on the overlapping region of MIRI MRS channel 1-short and the NIRSpec
IFU. The positional offset between these two instruments was $\Delta$RA =
0.37″and $\Delta$Dec = 0.67″. We also refined the position of the MIRI MRS
field by aligning it with Spitzer and Gaia data of the stellar sources
detected in the MIRI parallel image. Thus we aligned the NIRSpec and MIRI
spectral images before carrying out our analysis. For further elaboration and
in-depth information concerning this positional offset see Narang et al., (in
prep) and Federman et al. (2023b).

Figure 6: (left) The [Fe II] jet at 5.34 $\mu$m shown in blue overlaid with a
red rectangle encompassing the region used for making the PV diagram. (right)
The PV diagram for the [Fe II] line at 5.34 $\mu$m. The white circles
represent the center of the Gaussian fit to the velocity channels, while the
error bars represent the FWHM of the Gaussian fit. The white bar in the top
left corner is the PSF of the JWST MIRI at 5.34 $\mu$m, which is 39.2 au
(0.28″).

##  3 Jet from IRAS 16253-2429

###  3.1 Morphology

The observations of IRAS 16253-2429 carried out with the JWST NIRSpec IFU and
MIRI MRS instruments provide valuable insights about the protostar, revealing
atomic jets, molecular outflows, and envelope cavity structures at spatial
scales down to 30 au. In Figure 1, we present a comparison between the Spitzer
IRAC (IRAC 1, 2, and 4; Figure 1 a) observations of the protostar, covering a
field of view of 2′$\times$ 2′, and the NIRSpec IFU + MIRI MRS observations
(Figure 1 b) focused within the inner 7″ (also see Federman et al., 2023b).
Figure 1 (a) shows the full spatial extent of the protostellar outflow from
IRAS 16253-2429 as seen by Spitzer. The black rectangle shows the field of
view of the JWST observations. Figure 1(b) shows the zoomed-in inner $\sim$6″
region of the protostar as observed with JWST. The 4.17 $\mu$m continuum
emission is primarily the scattered light from the central source and is shown
in red; it traces the extent of the outflow cavity carved out of the envelope.
The molecular outflow traced by H2 0-0 S(11) emission at 4.18 $\mu$m is shown
in green. The atomic jet seen in the [Fe II] at 5.34 $\mu$m line is shown in
blue. The JWST observations reveal that IRAS 16253-2429 exhibits the
distinctive “wasp-waist” shape (Barsony et al., 2010), even within the inner
6″ region surrounding the protostar. Figure 1(b) further shows that the
scattered light cavity is much broader than the molecular emission from H2
(also see Federman et al. 2023b and Narang et al. in prep).

The wavelength range covered by the MIRI MRS and NIRSpec IFU instruments
encompasses several fine structure lines, including multiple lines of [Fe II].
In Figure 2, we present the line maps of all the detected [Fe II] lines
observed towards IRAS 16253-2429. These [Fe II] lines trace a highly
collimated jet originating from the central protostar. This is the first time
(at any wavelength) that a collimated jet has been detected from IRAS
16253-2429. The jet appears brighter on the northern side and weaker to the
south; but, this is because of the relatively higher line-of-sight extinction
toward the southern cavity (see Section 3.4).

We further find that there are knots at both ends of the jet (see Figure 2) at
around 400 au from the central protostar (also see Federman et al. 2023b). The
southern knot is best detected in the [Fe II] line at 5.34 $\mu$m, with a
tentative detection in the [Fe II] lines at 17.936 $\mu$m and 25.988 $\mu$m.
The northern knot is clearly seen in five of the six [Fe II] line maps
presented in Figure 2, although it is not fully mapped in the 5.34 $\mu$m
line, and is only faintly seen in the 24.5 $\mu$m line. If we draw a line from
the central protostar to the northern- and southern-most knots we find that
the jet seen in [Fe II] at 5.34 $\mu$m does not fall on this line. The
northern knot is clearly offset to the east from the axis of the rest of the
northern jet. We also find that the jet has a slight curvature, suggesting
that it might be precessing (also see Barsony et al. 2010).

In Figure 3, we present maps of other atomic fine-structure (FS) lines along
with the Br-$\alpha$ line at 4.052 $\mu$m. We find that the FS lines of [Ne
II] at 12.81 $\mu$m and the [Ni II] at 6.64 $\mu$m also trace the jet seen in
[Fe II]. The [Ne II] jet also shows faint knots at both ends similar to the
[Fe II] jets. The [S I] emission at 25.25 $\mu$m, however, is concentrated on
the protostar and does not appear to trace the jet (see Figure 3). The
Br-$\alpha$ line observed with NIRSpec appears to trace the jet and shows the
knot on the northern side (see Figure 3). Interestingly, Br-$\alpha$ emission
is stronger in the extended jets than at the central source. This suggests
that the H I lines observed from protostars via spatially unresolved
spectroscopy might have a large contribution due to the jet component; the
dominant jet component makes it difficult to isolate emission in this line
from accretion flows (see Hartmann et al., 2016; Federman et al., 2023b).

###  3.2 Jet width

Given the high signal-to-noise and high angular resolution at 5.34 $\mu$m
($\sim$ 0.28″ Law et al. 2023), it is possible to measure the width of the jet
as a function of distance from the protostar. The jet has a position angle of
23°(measured from North towards East). We started by taking slices aligned
perpendicular to the jet and averaged them over 3 pixels length along the jet
as shown in Figure 4(a). We then fit each of these averaged slices with a
Gaussian and the FWHM of the Gaussian is taken as the width of the jet in that
slice.

In Figure 4(b) we show the width of the jet (FWHM) as a function of distance
from the protostar. As can be seen from Figure 4(b), the protostellar jet
appears to expand as we move away from the central protostar. We measure the
FWHM of the slices adjacent to the protostar on either side and find the width
of the jet close to the central protostar to be $\sim$ 0.33 $\pm$ 0.03″. This
is comparable to the average PSF FWHM of the MIRI MRS at 5.34 $\mu$m of $\sim$
0.28″(Law et al., 2023).

In Figure 4(c) we show the deconvolved jet widths (obtained by subtracting out
the MIRI MRS PSF FWHM in quadrature) as a function of distance from the
central protostar. From a linear fit to the jet width as a function of
distance from the central protostar, we obtain the opening angle $\Theta$
($={\rm slope/2}$) as well as the jet width at the central position (zero
intercept). We find that the average slope of the fit is 0.09 $\pm$ 0.02,
corresponding to an opening angle of 2.6 $\pm$ 0.5°. The jet width
extrapolated to the source position is 23 $\pm$ 4 au, which we take as an
upper limit on the jet width at the source location. Federman et al. (2023b)
used the NIRSpec observation of the [Fe II] line at 4.88 $\mu$m to measure the
deconvolved width of the jet from IRAS 16253-2429 at its brightest position.
They found a deconvolved jet width to be 0.144″ which corresponds to $\sim
20.2$ au

Figure 7: The PV diagrams of the observed [Fe II] lines aligned to a common
average velocity.

###  3.3 Velocity structure in the jet

In order to explore the velocity structures within the jet, we first
constructed a moment 1 map of the [Fe II] line at 5.34 $\mu$m as shown in
Figure 5. The moment 1 map gives us the intensity-weighted velocity and is an
excellent diagnostic tool for understanding the kinematic structures within
the jets. As can be seen from Figure 5, the northern part of the jet is blue-
shifted with respect to the average velocity, while the southern jet is red-
shifted. We note that the velocity of the jet on both sides is fairly constant
and the maximum blueshift and redshift velocity difference is $\sim 130$ km
s-1. This is much larger than the systemic cloud velocity of $\sim$ 4 km s-1
(Aso et al., 2023). We assume the latter as VLSR.

To further investigate the velocity structure of the jet, we constructed a
Position-Velocity (PV) diagram, specifically for the [Fe II] line at 5.34
$\mu$m. In order to capture the full range of velocities, we carefully
selected a rectangular region around the protostar that is both wide and long
enough to encompass the entire jet, as illustrated in Figure 6 (a). We show
the PV diagram for the [Fe II] line at 5.34 $\mu$m in Figure 6 (b). From the
PV diagram, it is apparent that as we move away from the central protostar
(the driving source of the jet), the velocities on either sides of the jet
remain more or less constant. This is similar to what we found with the moment
1 map of [Fe II] at 5.34 $\mu$m (see Figure 5) as well. However, within the
inner region covering approximately 50 au, we witness rapid changes in
velocity.

We also constructed PV diagrams for the various other [Fe II] lines. We
computed the velocity of each slice of the spectral cube by converting the
wavelength shift with respect to the lab wavelengths (Koo et al., 2016; Tayal
& Zatsarinny, 2018) to a velocity shift. In Figure 7, we show PV diagrams of
all the [Fe II] lines that we have detected from IRAS 16253-2429. Initially,
the PV diagrams of the various [Fe II] lines did not align well since
uncertainties in the wavelength calibration can introduce offsets of as large
as 40 km s-1 even within a single channel (Argyriou et al., 2023). In Figure 7
we align the velocities in the PV diagrams by setting the mean velocity of the
jet seen in each line to zero. It can be seen from Figure 7 that the jets
display approximately symmetric velocities with respect to the central
protostar. The difference in the peak positive and negative velocities for all
the [Fe II] jets are similar, with an average value of 148 $\pm$ 12 km s-1.
This consistency in peak jet velocities seen in different [Fe II] lines
reinforces the notion of a coherent jet structure originating from the central
protostar.

Using the PV diagrams we compute the jet velocity as half of the average peak-
to-peak velocity. The average peak-to-peak velocity is 148 $\pm$ 12 km s-1,
which gives the radial jet velocity as 74 $\pm$ 6 km s-1. After correcting for
the inclination of the jet using the measured inclination angle of 64.1 $\pm$
0.5° from Aso et al. (2023), we obtain the peak velocity of the jet to be 169
$\pm$ 15 km s-1.

###  3.4 Measuring extinction along the jet

To derive the physical properties of the jet from the observed line maps, we
first measure and quantify the line of sight extinction. To determine the
spatially varying extinction towards IRAS 16253-2429, we used the molecular H2
0-0 S lines. Molecular H2 lines have previously been used with Spitzer IRS to
determine the extinction values (e.g., Barsony et al., 2010). If we assume
that the H2 is in local thermodynamical equilibrium (LTE) with an otho-para
ratio of 3 we can construct the rotation diagram for the H2 lines (Figure 9
and Figure 11).

In this diagram, the natural logarithm of the column density of H2 molecules
in the $J^{th}$ rotational state, $N_{u}$, divided by the degeneracy of that
state ($g_{u}$ is the rotational degeneracy equal to 2J + 1 and $g_{s}$ is the
spin degeneracy), ln$(N_{u}/g_{u}g_{s})$, is plotted as a function of the
upper state energy, $E_{u}$ (energy of the $J^{th}$ rotational state) (see,
e.g., Neufeld et al., 2009; Maret et al., 2009; Barsony et al., 2010; Manoj et
al., 2013). Assuming that H2 emission is optically thin, then the number of H2
molecules in the $J^{th}$ rotational state is given by:

$N_{u}=\frac{4~{}\pi~{}F_{J,J-1}~{}\lambda_{J,J-1}}{h~{}c~{}A_{J,J-1}}$ (1)

where $h$, $c$, $F_{J,J-1}$, $\lambda_{J,J-1}$ and $A_{J,J-1}$ are the
Planck’s constant, speed of light in vacuum, line flux within the aperture,
wavelength, and Einstein A-coefficient, respectively, corresponding to the
transition $J\rightarrow J-1$. We can compute an average rotational
temperature $T_{rot}$, for a relatively small range in $E_{u}$ by fitting a
straight line to the points in the observed rotational diagram. The negative
of the inverse of the slope gives us the $T_{rot}$.

Figure 8: (left) The 6 apertures along the jet where we measure the extinction
and jet properties. The center of aperture 3 (blue circle) marks the position
of the protostar. (right) The measured extinction value along the jet.

Under the LTE assumption, any deviation from this straight line in the narrow
range of $E_{u}$ is due to extinction. By correcting for these deviations we
derive the line of sight extinction towards the protostars. To do so, we can
simultaneously fit the rotation diagram and the correction for extinction. We
restrict our analysis to H2 lines between 6-18 $\mu$m (from H2 0-0 S(1) to
S(6)) to minimize non-LTE effects and also to avoid contamination from
molecular CO (also see Rubinstein et al., in prep). We used the extinction law
from Klaus Pontoppidan (hereafter KP v05; private communication) presented in
(Chapman et al., 2009) to derive the extinction towards the protostar (also
see section 3.5). We select apertures of radius 0.8″ spaced 1.1″ apart along
the jet (see Figure 8a) and measure the line fluxes within them. We then fit
the extinction value $A_{V}$ and two rotation temperatures simultaneously, one
temperature representing cooler gas between H2 0-0 S(1) and S(4) and the other
slightly hotter gas between H2 0-0 S(3) and S(6)) to the H2 0-0 S lines. In
Figure 9 we show the H2 rotation diagram for aperture 3 (centered on the
protostar). We can clearly see how correcting for extinction, especially for
the H2 S(3) line leads to an improved fit to the temperatures. Figure 11 in
the Appendix shows the rest of the H2 rotation diagrams. The rotation
temperature and $A_{V}$ values for different apertures are listed in Table 2.

Figure 9: The H2 rotation diagram for aperture 3 (centered on the protostar)
in blue. Also shown are the observed points (red) as well as the fit to the
dereddened points (blue solid and dashed lines). The extinction value $A_{V}$
as well as the rotation temperatures are listed in the Figure.

In Figure 8(b), we show the variation of the line of sight extinction along
the length of the jet. We find that the extinction rapidly rises as we move
towards the protostar from North to South, attaining a maximum at the
protostellar position, after which it remains relatively constant. The details
of the H2 excitation analysis will be provided in a follow-up paper.

Figure 10: The MAPPINGS shock model grid (nomogram) with the extinction
corrected ratios for the fine structure lines of [Fe II] and [Ne II]. The
ratio of [Ne II] at 12.81 $\mu$m and [Fe II] at 5.34 $\mu$m in blue runs along
the y-axis and is sensitive to the shock velocity while the ratio of [Fe II]
lines at 17.93 and 5.34 $\mu$m in green that almost runs parallel to the
x-axis is sensitive to the preshock density. The dashed line is an upper limit
to the log ([Fe II] 17.936 $\mu$m/ 5.34 $\mu$m).

###  3.5 Physical conditions along the jet

We next examined the physical conditions (shock velocity and density) along
the jet. We used the 1D shock and photoionization code MAPPINGS V (Sutherland
et al., 2018) to model the shock. Furthermore, we used the [Fe II] collisional
excitation network from (Tayal & Zatsarinny, 2018) instead of the CHIANTI
defaults. Our shock models were computed for solar abundance (Asplund et al.,
2009) with a frozen-in magnetic field $B=15.8\,\mu
G\times\sqrt{n/1000\,cm^{-3}}$. We implement the full magnetic and UV
radiative precursor treatment by iterating and integrating the shock model to
convergence.

We chose to use the fine structure lines of [Fe II] and [Ne II] as these lines
have the largest extent in the jet. In Figure 10 we have plotted the MAPPINGS
shock model grids (nomogram). The predicted ratio of [Ne II] at 12.81 $\mu$m
and [Fe II] at 5.34 $\mu$m in blue runs along the y-axis and is sensitive to
the shock velocity ($v_{s}$) while the predicted ratio of [Fe II] lines at
17.93 and 5.34 $\mu$m in green that almost runs along the x-axis is sensitive
to the preshock density, ($\rho_{s}$). By placing the observed line ratios
onto this nomogram we can directly read off the shock speed and the preshock
density.

We measured the line fluxes along the jet using the same aperture as Figure 8
(left). We placed the dereddened flux ratios for the [Fe II] and [Ne II] lines
onto the grid in Figure 10 as colored symbols (the color corresponds to the
aperture in Figure 8 (left)) and read off the shock speed and preshock
density. We have listed the shock speed and preshock density in Table 2. We
find that away from the central protostar the preshock density rapidly falls,
while the shock speed remains roughly constant.

The mass loss rate from the protostar $\dot{M}_{out}$ can be expressed in
terms of the pre-shock density, velocity of the gas entering the shock, and
jet radius as:

$\dot{M}_{out}=\pi\times R^{2}\times\rho_{s}v\times m_{\mathrm{H}}$ (2)

where $R$ is the radius of the jet, $\rho_{s}$ is the pre-shock density, $v$
is the velocity of the gas entering the shock, and $m_{\mathrm{H}}$ is the
mass of the hydrogen atom. We have also measured the velocity of the shocked
gas ($v_{jet}$) in internal working surfaces using the PV diagrams. Based on
different jet models the velocity of gas $v$ can either be the bulk velocity
of the jet as derived from the PV diagram ($v=v_{jet}$) (e.g., pulsating jet
models of Raga et al., 1990; Raga & Kofman, 1992) or the sum of the shock
velocity and the jet velocity ($v=v_{jet}+v_{s}$) (e.g., Rubinstein et al.,
2023).

Using the pre-shock density and shock speed derived from the shock grid as
well as the jet velocity derived from the PV diagram, and assuming that the
jet has a cylindrical cross-section with radius equal to half the jet width
measured at the location, we have computed the mass loss rate along the jet
from the protostar. The mass loss rate for the various apertures is listed in
Table 2. We find that the upper limit to the mass loss rate $\dot{M}_{out}$
from the central protostar is 1.1 $\times 10^{-10}$ $M_{\odot}$/yr.

In computing the mass loss rate $\dot{M}_{out}$ from the protostar, the
largest uncertainty is not due to any measurement uncertainties, but due to
the uncertainty in the adopted extinction law. Chapman et al. (2009) have
shown that the extinction law measured towards molecular clouds is highly
variable and not only depends upon the molecular cloud but also the
environment. To quantify the effects of using different extinction laws on our
results, we repeated our analysis of computing extinction and estimating the
mass loss rate using two additional extinction laws, that of McClure (2009)
and Hensley & Draine (2021). We compute the mass loss rate from the central
protostar at the base of the jet (from aperture 3) assuming $v=v_{jet}+v_{s}$.
We find that the mass loss rate obtained using the extinction law from McClure
(2009) is 4.2 $\times 10^{-11}~{}M_{\odot}$ yr -1 and using Hensley & Draine
(2021) is 3.2 $\times 10^{-11}~{}M_{\odot}$ yr -1. These $\dot{M}_{out}$
values and the ones listed in Table 2 (for aperture 3) are within a factor of
3.5, indicating that the upper limit that we obtain for mass loss rate,
$\dot{M}_{out}~{}\leq~{}1.1\times 10^{-10}$ is robust.

##  4 Discussion

Our JWST observations reveal a highly collimated atomic jet in the IR fine
structure lines and H I lines from the low luminosity and very low mass
protostar IRAS 16253-2429. No collimated molecular jet is detected, and
although IR H2 emission is detected along the outflow cavity, it has much
wider opening angle and does not spatially follow the collimated atomic jet.
The jet seen in the FS lines appears brighter on the blueshifted north-eastern
side that is facing us, and is fainter to the south-west; this is due to the
higher line-of-sight extinction toward the south-western cavity (see from
Figure 8).

The jets on either side appear to have similar velocities that more or less
remain constant as a function of distance from the driving source. The jet
velocity, corrected for inclination effects (using the inclination of 64.1°),
is 169 $\pm$ 15 km s-1. The collimated jet seen in the JWST line images
extends up to about 400 au (see Figure 7; 916 au after correcting for
inclination) from the driving source on either side. This implies a dynamical
timescale of $\sim$26 yr, suggesting that the jet seen in the JWST IFUs was
ejected only in the last 26 yr.

The width of the jet (FWHM) at the base is marginally resolved and is $\sim$
23 $\pm$ 4 au, and the jet tends to widen as it moves away from the central
source with an opening angle $\Theta$ = 2.6°. If the jet is ballistically
confined, then the opening angle can be expressed as a ratio of the sound
speed within the jet and the jet speed.

$\Theta=\frac{c_{s}}{v_{\mathrm{jet}}}$ (3)

where $v_{\mathrm{jet}}$ is the jet velocity of 169 $\pm$ 15 km s-1 and
$c_{s}$ is the sound speed. If we assume that the shock temperature is 10,000
K (e.g., Ray & Ferreira, 2021) then the sound speed $c_{s}$ = 10 km s-1, which
gives an opening angle of 3.4 $\pm$ 0.3°. This is roughly consistent with the
opening angle we measure from the observed jet widths. This suggests that the
jet is likely ballistically confined, at least within the central $\lesssim$
1000 au of the driving source; other confinement mechanisms, however, could be
at play at larger distances.

The physical conditions of the jet derived from our analysis indicate shock
velocities of 42$-$54 km s-1 from the base to one end of the jet, suggesting
that there is no significant acceleration along the jet. The preshock
densities estimated are in the range of $\sim$ 10-2000 cm-3, indicating that
the jet is rather tenuous.

Table 2: The extinction $A_{V}$, the rotation temperature of the H2 gas (Trot1 and Trot2), the shock velocities, preshock densities, jet radius $R$ (computed using the opening angle $\Theta$ and the jet width at the base) along with the mass outflow rate $\dot{M}_{out}$ for each of the apertures shown in Figure 8. 222∗ Aperture 1 is a lower limit due to the fact that the northern knot is not entirely covered by the MIRI channel-1 FOV. Aperture | Distance from | $A_{V}$ | Trot1 | Trot2 | $v_{s}$ | $\rho_{s}$ | $R$ |  $\dot{M}_{out}$ for |  $\dot{M}_{out}$ for
---|---|---|---|---|---|---|---|---|---
| Protostar |  |  |  |  |  |  | $v=v_{jet}+v_{s}$ | $v=v_{jet}$
| (″) |  | (K) | (K) | (km s-1) | (cm-3) | (au) | ($M_{\odot}$ yr -1) | ($M_{\odot}$ yr -1)
1∗ | -2.2 | 12.3 $\pm$ 0.6 | 526 $\pm$ 34 | 794 $\pm$ 65 |  $>$42 |  $>$6$\times 10^{1}$ | 25.3 |  $>$1.5 $\times 10^{-11}$ |  $>$ 1.2$\times 10^{-11}$
2 | -1.1 | 16.2 $\pm$0.4 | 585 $\pm$ 41 | 873 $\pm$ 87 | 51 | 2$\times 10^{2}$ | 18.4 | 2.8$\times 10^{-11}$ | 2.1$\times 10^{-11}$
3 | 0 | 26.6 $\pm$ 0.5 | 661 $\pm$ 72 | 847 $\pm$ 103 | 54 | 2$\times 10^{3}$ | 11.5 | 1.1$\times 10^{-10}$ | 8$\times 10^{-11}$
4 | 1.1 | 26.1 $\pm$ 0.4 | 685 $\pm$ 60 | 852 $\pm$ 93 | 52 | 3$\times 10^{2}$ | 18.4 | 4.2$\times 10^{-11}$ | 3.2$\times 10^{-11}$
5 | 2.2 | 23.9 $\pm$ 0.4 | 622 $\pm$ 55 | 840 $\pm$ 83 | 54 | 1$\times 10^{1}$ | 25.3 | 2.7$\times 10^{-12}$ | 2$\times 10^{-12}$
6 | 3.3 | 24.2 $\pm$0.4 | 672 $\pm$ 51 | 882 $\pm$ 108 | 52 |  $<$ 2 | 32.3 |  $<$8.6$\times 10^{-13}$ |  $<$6.6$\times 10^{-13}$

The mass loss rate measured from the central protostar is
$\dot{M}_{out}~{}\leq$ 1.1 $\times 10^{-10}~{}M_{\odot}$ yr -1. Even when
different extinction laws are used, the mass loss rate changes only by a
factor of 3.5 or so. The estimated mass loss rate appears low, even for a very
low mass protostar. The simultaneous accretion rate derived from the OH lines
in the MIRI spectrum of IRAS 16253-2429 is $\dot{M}_{acc}~{}\sim 1.3\times
10^{-9}~{}M_{\odot}$ yr -1 (Watson et al. in prep), commensurate with the low
mass loss rate observed. The protostar appears to be going through a phase of
low mass accretion/ejection currently. It is, however, driving a highly
collimated jet, even while in the relatively quiet phase, suggesting that
collimated jets are present even when protostars are passing through a low
accretion phase (also see Aso et al., 2023).

Although the protostar IRAS 16253-2429 is currently in a quiescent phase, it
must have gone through earlier epochs of high mass accretion/ejection, and is
also very likely to go through active high accretion phase later in its life,
for the following reasons.

The current mass of the central source based on the ALMA observations of the
Keplerian disk is about 0.15 $M_{\odot}$ (Aso et al., 2023). At the current
mass accretion rate, it would take more than 100 Myr to accrete a mass of 0.15
$M_{\odot}$. IRAS 16253-2429 is a Class 0 protostar, which is likely to be
$\ll 1$ Myr old (e.g., Dunham et al., 2014), suggesting that the protostar
must have had a period of much higher accretion rate at an earlier stage
during which most of the current mass was accreted. Additionally, wide
cavities are already present in the envelope of IRAS 16253-2429 (see Federman
et al., 2023b), which must have been carved out by powerful jets and outflows
associated with an earlier high accretion phase. All these evidence indicate
that IRAS 16253-2429 must have been accreting at a significantly higher rate
at an earlier time in its life. Zakri et al. (2022) have shown that during
these high accretion phases, the protostar can accrete as much as 100% of its
current mass.

The envelope mass of IRAS 16253-2429 estimated from mm observations is in the
range of 0.2 to 1 $M_{\odot}$ (Stanke et al., 2006; Enoch et al., 2008; Tobin
et al., 2011) depending on the size of the core considered, the opacity values
used, and the assumed dust temperature. All measurements and our modeling
efforts are consistent with a mass of $\gg$ 0.2 $M_{\odot}$ within an envelope
radius of $\sim$ 4200 au, suggesting that there is significantly more mass in
the envelope of IRAS 16253-2429 than in the central object. Given the
available mass in the envelope, it is reasonable to assume that at least an
additional 0.15 $M_{\odot}$ will be accreted onto the central source. This
would make the final mass of IRAS 16253-2429 to be $\sim$0.3$~{}M_{\odot}$,
which is where the stellar initial mass function peaks (Bastian et al., 2010).

At the current accretion rate, it would take $>100$ Myr to accrete this mass,
strongly suggesting that the protostar will again go into a high mass
accretion/ejection phase. This is consistent with the results of Zakri et al.
(2022), who find that Class 0 protostars undergo bursts every few hundred
years that last only a few decades; hence it is more likely that a protostar
is observed in between bursts. In the case of one Class 0 protostar, HOPS 383,
the pre and post- burst periods are characterized by a quiescence with a very
low accretion rate (Zakri et al., 2022). Our JWST observations appear to have
captured IRAS 16253-2429 during the interlude between the accretion bursts,
i.e. in the quiescent phase, and yet driving a tenuous but highly collimated
jet.

The low mass loss rate that we derive for the protostar is also consistent
with the non-detection of the molecular jet from the source. Tabone et al.
(2020) have argued that self-shielding of the jet from FUV photons generated
from the central source or from shocks along the jet requires much higher mass
loss rates from the protostar.

##  5 Summary and Conclusions

In this work, we report the discovery of a highly collimated jet in multiple
[Fe II], other fine structure lines, an H I originating from the Class 0
protostar IRAS 16253-2429, the least luminous (and least massive) source in
the Investigating Protostellar Accretion (IPA) program. No collimated
molecular jet is detected in CO, SiO or H2 in the IR with JWST and in the
(sub-)mm with ALMA. We determine the following properties of the
jets/outflows:

  1. 1.

The jet (as seen in [Fe II] line at 5.34 $\mu$m) is 23 $\pm$ 4 au wide at the
base with an opening angle of 2.6 $\pm$ 0.5°.

  2. 2.

Based on our analysis of the PV diagram of various [Fe II] lines, we find that
the jet has a velocity of 169 $\pm$ 15 km s-1 (after correcting for
inclination). The jets on either side of the protostar appear to have similar
velocities that remain almost constant as a function of distance from the
center.

  3. 3.

We determine the line of sight extinction $A_{V}$ along the jet using H2 lines
in the cavity, which gives $A_{\mathrm{V}}$ values ranging from $\sim$ 12 to
$\sim$27.

  4. 4.

Using the extinction corrected flux ratios of [Ne II] at 12.81 $\mu$m and the
[Fe II] lines at 5.34 $\mu$m and 17.936 $\mu$m, we derive a shock velocity of
54 km s-1 and a preshock density of 2.0$\times 10^{3}$ cm-3 at the base of the
jet.

  5. 5.

Assuming a cylindrical jet with a cross-sectional width of 23 au and using the
jet velocity of 169 km s-1 along with the derived preshock density and shock
velocity, we compute the mass loss rate from the jet to be between 0.32-1.1
$\times 10^{-10}~{}M_{\odot}$ yr -1.

The low mass loss rate derived is consistent with simultaneous measurements of
low accretion rate of $\lessapprox 10^{-9}~{}M_{\odot}$ yr -1 for IRAS
16253-2429 (Watson et al. in prep), indicating that the protostar is likely in
the quiescent accretion phase. Our results show that even a very low mass
protostar in a quiescent phase with low accretion rate can still drive a
highly collimated jet.

##  6 Acknowledgment

This work is based on observations made with the NASA/ESA/CSA James Webb Space
Telescope. The data were obtained from the Mikulski Archive for Space
Telescopes at the Space Telescope Science Institute, which is operated by the
Association of Universities for Research in Astronomy, Inc., under NASA
contract NAS 5-03127 for JWST. These observations are associated with program
#1802. Support for SF, AER, STM, RG, WF, JG, JJT and DW in program #1802 was
provided by NASA through a grant from the Space Telescope Science Institute,
which is operated by the Association of Universities for Research in
Astronomy, Inc., under NASA contract NAS 5-03127. A.C.G. has been supported by
PRIN-MUR 2022 20228JPA3A “The path to star and planet formation in the JWST
era (PATH)” and by INAF-GoG 2022 “NIR-dark Accretion Outbursts in Massive
Young stellar objects (NAOMY)”. G.A. and M.O., acknowledge financial support
from grants PID2020-114461GB-I00 and CEX2021-001131-S, funded by
MCIN/AEI/10.13039/501100011033.

JWST (NIRSpec, MIRI)

##  Appendix A Rotation diagram along the jet

Figure 11: The H2 rotation diagram for all aperture as shown in Figure 8. Also
shown are the observed points (red) as well as the fit to the dereddened
points (blue solid and dashed lines). The extinction value $A_{V}$ as well as
the rotation temperatures are listed in the figure.

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