11institutetext: Maynooth University Department of Experimental Physics,
National University of Ireland Maynooth, Maynooth Co. Kildare, Ireland
22institutetext: European Southern Observatory, Karl-Schwarzschild-Strasse 2,
85748, Garching bei München, Germany 33institutetext: Johns Hopkins
University, 3400 N. Charles Street, Baltimore, MD 21218,USA 44institutetext:
Space Telescope Science Institute, 3700 San Martin Dr, Baltimore, MD 21218,
USA 55institutetext: Centre for Extragalactic Astronomy, Department of
Physics, Durham University, South Road, Durham DH1 3LE, UK 66institutetext:
Institute for Computational Cosmology, Department of Physics, University of
Durham, South Road, Durham DH1 3LE, UK 77institutetext: Dipartimento di
Fisica, Universitá degli Studi di Milano, via Celoria 16, Milano, Italy
88institutetext: Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112,
21029 Hamburg, Germany 99institutetext: Instituto de Astrofísica e Ciências do
Espaco, Universidade de Lisboa, OAL, Tapada da Ajuda, P-1349-018 Lisboa,
Portugal

# A spectacular jet from the bright 244-440 Orion proplyd

The MUSE NFM view ††thanks: Based on observations obtained with the MUSE
spectrograph on the Very Large Telescope on Cerro Paranal (Chile), operated by
the European Southern Observatory (ESO). Program ID: 0104.C-0963(A)

A. Kirwan 11 C. F. Manara 22 E. T. Whelan 11 M. Robberto 3344 A. F. McLeod
5566 S. Facchini 77 G. Beccari 22 A. Miotello 22 P. C. Schneider 88 A. Murphy
11 S. Vicente 99

In this work we present the highest spatial and spectral resolution integral
field observations to date of the bipolar jet from the Orion proplyd 244-440
using Multi-Unit Spectroscopic Explorer (MUSE) narrow-field mode (NFM)
observations on the Very Large Telescope (VLT). We observed a previously
unreported chain of six distinct knots in a roughly S-shaped pattern, and by
comparing them with Hubble Space Telescope (HST) images we estimated proper
motions in the redshifted knots of 9.5 mas yr-1 with an inclination angle of
$None$, though these quantities could not be measured for the blueshifted
lobe. Analysis of the [Fe ii] and [Ni ii] lines suggests jet densities on the
order of $\sim 10^{5}\leavevmode\nobreak\ {\rm cm}^{-3}$. We propose that the
observed S-shaped morphology originates from a jet launched by a smaller
source with $M_{\star}<0.2$ $M_{\odot}$ in orbital motion around a larger
companion of $M_{\star}\simeq 0.5$ $M_{\odot}$ at a separation of $30-40$ au.
The measured luminosities of the knots using the [O i]$\lambda 6300$ Å and [S
ii]$\lambda 6731$ Å lines were used to estimate a lower limit to the mass-loss
rate in the jet of $1.3\times 10^{-11}$ $M_{\odot}$ yr-1 and an upper limit of
$10^{-9}$ $M_{\odot}$ yr-1, which is typical for low-mass driving sources.
While the brightness asymmetry between the redshifted and blueshifted lobes is
consistent with external irradiation, further analysis of the [Ni ii] and [Fe
ii] lines suggests that photoionization of the jet is not likely to be a
dominant factor, and that the emission is dominated by collisional excitation.
The dynamical age of the jet compared to the anticipated survival time of the
proplyd demonstrates that photoevaporation of the proplyd occurred prior to
jet launching, and that this is still an active source. These two points
suggest that the envelope of the proplyd may shield the jet from the majority
of external radiation, and that photoionization of the proplyd does not appear
to impact the ability of a star to launch a jet.

###### Key Words.:

ISM: jets and outflows – ISM: individual: Orion Nebula – stars: pre-main-
sequence – stars: individual: 244-440 – protoplanetary disks

##  1 Introduction

Figure 1: HST (left) and MUSE (right) field images of the region around
244-440. The fields and insets in both images are shown in the [O i]$\lambda
6300$ Å emission. The stars $\theta^{2}$ Orionis A and B are seen just below
the proplyd to the southwest and southeast, respectively. The inset in the
MUSE panel is a flux-integrated image from our current NFM observations
discussed in Section 2.1.

Astrophysical jets are near-ubiquitous phenomena in the evolution of low- and
intermediate-mass young stellar objects (YSOs, Frank et al., 2014), and have
similarly been observed in high-mass YSOs as well (Marti et al., 1993; McLeod
et al., 2018). The formation of clumpy, shock-heated condensations known as
Herbig-Haro (HH) objects is linked with ejection events, and as such the study
of these objects allows us to better understand the mass-loss history of the
star. The morphology and physical conditions of the jets are determined by the
characteristics of the ejection mechanism as well as by the surrounding
environments: single, isolated stars may launch straight, bi-polar outflows,
while driving sources in denser regions may experience side winds from more
massive stars, possibly compounded with a high proper motion, that induce a
curvature in the jet (Bally & Reipurth, 2001; Raga et al., 2009); in binary
systems, companions may induce warping in the inner disk, precession, or
orbital motion that causes a “wiggling” jet morphology (Masciadri & Raga,
2002; Lai, 2003; Murphy et al., 2021; Erkal et al., 2021; Kirwan et al.,
2022).

The exact mechanism of jet launching is not known, but current research has
found a strong correlation between ejection events and mass accretion and
circumstellar disk dynamics (Ray & Ferreira, 2021; Whelan, 2014; Nisini et
al., 2018). Studying such ejections in a variety of environments is crucial to
fully understand the physics behind them. Regions such as the Orion Nebula
Cluster (ONC) provide a unique laboratory for studying the early stages of
star formation, due to the high stellar density and the relative closeness of
the region ($d\leavevmode\nobreak\ \sim$ 400 pc). Moreover, the environment
allows for jets to be examined under the most extreme conditions.

Many YSOs within the ONC have been found to possess externally illuminated,
photo-evaporating disks often surrounded by nebular structures and ionization
fronts (IFs). With the discovery of these “proplyds” (PROtoPLanetarY DiskS)
came the observations that many of them are associated with jets and HH
objects (O’Dell et al., 1997; Bally & Reipurth, 2001; Ricci et al., 2008).
They typically are one-sided and exhibit a C-shaped curvature pointed away
from the massive ionizing stars within the cluster. Proplyds in the dense,
inner regions of the ONC present us with a lifetime problem as the measured
mass-loss rates due to ultraviolet (UV) irradiation are too high and should
rapidly evaporate the disk (Clarke, 2007). However, near-infrared (NIR) excess
is still observed in $\sim 80$% of the ONC stars, implying that their disks
may survive longer than predicted. This implies that either initial disk
masses are large ($>1$ $M_{\odot}$) and therefore unstable, or that the
massive ionizing stars have ages $\leq 0.1$ Myr, which is short compared to
the region average ($2-4$ Myr, see Beccari et al., 2017). Estimates of disk
masses in the inner regions of the ONC have not generally been larger than
$10^{-2}$ $M_{\odot}$ (Henney & O’Dell, 1999), and more recent surveys with
the Atacama Large Millimeter/submillimeter Array (ALMA) have indicated a
maximum dust mass of $\sim 80\leavevmode\nobreak\ M_{\oplus}$ (Eisner et al.,
2018). A combination of age spread and stellar dynamics, with the youngest
stars migrating inward to the central regions, may help solve the lifetime
discrepancy (Winter et al., 2019). This should have direct implications on the
morphology of the jets and on the accretion-outflow connection in
photoevaporated disks, granting deeper insights into the mechanics of jet
launching and the impact of the environment.

The giant proplyd 244-440, known also as V* V2423 Ori and HH 524, is located
at a distance of $400\pm 23$ pc (Gaia Collaboration et al., 2021). With a
width of 3$\aas@@fstack{\prime\prime}$5 ($\sim 1400$ au), it is one of the
largest proplyds in the ONC (Bally, O’Dell, & McCaughrean, 2000, hereafter
BOM2000). It is not located in the inner core of the ONC, as its large IF of
radius $\sim 2\aas@@fstack{\prime\prime}0$ ($\sim 800$ au) points toward
$\theta^{2}$ Orionis A, located southwest of the proplyd and beyond the
prominent photodissociation region (PDR) known as the “Orion Bar,” as shown in
Figure 1. Hubble Space Telescope (HST) images in H$\alpha$ were interpreted by
Bally & Reipurth (2001) as showing a nearly edge-on disk in silhouette with a
size of approximately 0$\aas@@fstack{\prime\prime}$15 $\times$
0$\aas@@fstack{\prime\prime}$6, and a tilted jet with a $\sim$$$ difference
between the jet axis and the disk minor axis. BOM2000 further note that the
star appears offset from the center of the disk by $\sim$
0$\aas@@fstack{\prime\prime}$1, and suggest a binary system where one star is
hidden within the disk. Spectral observations of the source suggest that it is
a low-mass star (¡ 1 $M_{\odot}$; see Appendix A).

Henney & O’Dell (1999) initially proposed a mass-loss rate for the proplyd of
$\sim 1.5\times 10^{-6}$ $M_{\odot}$ yr-1, although Winter et al. (2019) posit
that this is an overestimation, suggesting instead a mass-loss rate on the
order of $\sim 5\times 10^{-8}$ $M_{\odot}$ yr-1. Direct estimates of the
disk$+$envelope mass are difficult, but range from as high as 0.01 $M_{\odot}$
(Bally et al., 1998) based on millimeter measurements, to as low as $\sim
5\times 10^{-3}$ $M_{\odot}$ with the VLA (Sheehan et al., 2016).111We note
that the VLA is free$-$free (FF) emission dominated, and while it is possible
that the proplyds are FF dominated as well, this latter value is presented
with caution. While the values of Bally et al. (1998) suggest evaporation
times $t_{e}\sim 10^{4}$ yr (see also Henney & O’Dell, 1999), those of Sheehan
et al. (2016) and Winter et al. (2019) suggest $t_{e}\sim 0.1-0.2$ Myr, which
ultimately set a lower limit on the age of the proplyd.

In this paper, we present the first detailed analysis of the bi-polar jet
associated with this proplyd. In Section 2 we discuss our MUSE data and data
reduction process. In Section 3, we outline a proper motions study and offer
an explanation of the complex curvature of the jet. The extinction, physical
conditions, and mass-loss rate are estimated in Section 4. We discuss our
results in Section 5 and provide a summary in Section 6.

##  2 Observations & data reduction

###  2.1 VLT/MUSE

Observations of proplyd 244-440 were obtained on 23 October 2019 with the
Multi-Unite Spectroscopic Explorer (MUSE) under program ID 104.C-096 (PI: C.
F. Manara). The instrument was operated in narrow field mode (NFM) with
adaptive optics (AO) under clear sky conditions, which allows sampling of the
target with a spatial resolution of 0$\aas@@fstack{\prime\prime}$025 pixel-1.
The image quality delivered by the AO system was measured in the data as $\sim
0\aas@@fstack{\prime\prime}14$. For at least 50% of the total observation, the
coherence time was $>6$ ms. Data reduction was performed with the MUSE
pipeline (v2.8) with ESO Reflex using standard calibrations and recipes (for a
detailed description of the MUSE data reduction pipeline, see Weilbacher et
al., 2020). Further details of the data reduction process will be described in
a forthcoming paper. A final cube was produced spanning the entire nominal
wavelength range of the MUSE instrument ($\sim 4750-9350$Å), with a field of
view (FOV) of 8$\aas@@fstack{\prime\prime}$4 $\times$
8$\aas@@fstack{\prime\prime}$6.

This final cube was further processed in Python to remove local continuum and
nebular contributions, using a method similar to that described by Agra-
Amboage et al. (2009). For each spectral emission region of interest, a
subcube spanning $\sim 100$Å was extracted from the primary cube and
continuum-subtracted by fitting a second-order polynomial to the spectrum at
each spaxel in the subcube (see Kirwan et al., 2022, for a description of this
process and considerations of the wavelength-dependence of the PSF).
Additionally, to account for nebular background and foreground line emission
we constructed a mean, local background spectrum for each subcube by sampling
regions away from the proplyd envelope but close enough to be representative
of the emission profile, and subtracted this from the subcube. Extinction was
calculated using the Balmer decrement under Case B assumptions, which is
further discussed in Section 4.1.

###  2.2 Archival HST images

Two epochs of archival HST Wide-Field Planetary Camera (WFPC2) images were
used to estimate the proper motions of the knots (see Section 3.1). The
observations were obtained on 14 November 1995 and 17 September 1998 as part
of the General Observer programs GO 5976 (PI: J. Bally) and GO 6603 (PI: C.
O’Dell), respectively. We have focused on the F631N ([O i]$\lambda 6300$ Å,
both epochs) filters in this study, which have total exposure times $t_{\rm
exp}=2100$ s and $t_{\rm exp}=1200$ s respectively. Additionally, GO 5976 has
observations in the F673N ($t_{\rm exp}=2100$ s) and F791W filters ($t_{\rm
exp}=200$ s), and GO 6603 has observations in the F656N ($t_{\rm exp}=600$ s)
and F814W filters ($t_{\rm exp}=60$ s). The jet is seen in the F631N and F673N
filters, while only the envelope is seen in the F791W and F814W filters. A
description of the observations can be found in O’Dell et al. (1997).222The
science-ready data described here may be found on the MAST archive at
http://dx.doi.org/10.17909/y66h-8p10. To increase the S/N, we combined the [O
i] images for each epoch into two individual stacks. We then separately
aligned each stack of images to its own common reference frame in Python using
the mpdaf333https://github.com/musevlt/mpdaf function align_with_image, and
median-combined the two epochs separately. The jet-free images for each epoch
were finally subtracted from the [O i]$\lambda 6300$ Å images to better
isolate the jet emission, as shown in Figure 3.

Table 1: Emission line fluxes

$\lambda_{\mathrm{air}}$ (Å) | Ion | FE1 | FE2 | FE3 | FW3 | FW2 | FW1
---|---|---|---|---|---|---|---
5158.8 | [Fe iii] | … | … | … | … | 14.87 | 13.23
6300.3 | [O i] | 70.15 | 26.98 | 73.41 | 160.32 | 183.44 | 99.47
6363.8 | [O i] | 23.47 | 7.47 | 25.75 | 53.78 | 60.33 | 39.66
6716.4 | [S ii] | 15.48 | … | 15.95 | 52.92 | 71.16 | 49.25
6730.8 | [S ii] | 28.59 | … | 33.08 | 98.02 | 146.27 | 101.20
7155.2 | [Fe ii] | 11.88 | 8.38 | 11.03 | 20.39 | 19.78 | 17.19
7172.0 | [Fe ii] | 2.71 | 2.47 | 2.99 | 6.45 | 5.72 | 4.80
7255.8 | [Ni ii] | … | … | … | … | 7.89 | 4.27
7320.0 | [O ii] | 71.86 | 31.25 | 46.76 | … | … | …
7330.2 | [O ii] | 54.27 | 23.45 | 36.61 | … | … | …
7377.8 | [Ni ii] | 8.11 | 4.82 | 5.72 | 12.36 | 12.77 | 12.22
7388.2 | [Fe ii] | 2.10 | … | … | 3.35 | 3.20 | 2.80
7411.6 | [Ni ii] | 1.29 | … | … | … | 1.68 | 2.04
7452.5 | [Fe ii] | 2.83 | 2.23 | … | 5.78 | 5.81 | 5.65
8578.7 | [Cl ii] | 2.64 | … | … | 3.74 | 2.36 | 5.51
8617.0 | [Fe ii] | 13.64 | 9.27 | 12.27 | 23.93 | 23.47 | 20.48
9052.0 | [Fe ii] | … | … | … | 6.44 | 5.54 | 4.46
444Dereddened mission line fluxes for the knots in the Proplyd 244-440 jet, in
units of 10-17 erg/s/cm2. Line fluxes were computed through the Gaussian
fitting of spectral profiles extracted from the aperture sizes given in
Section 3.1. The average uncertainty in the lines fluxes is $\sim$5 $\times$
10-18 erg/s/cm2, which we calculate from the RMS noise in adjacent portions of
the spectrum. Empty entries denote positions where either no knot emission is
observed or the emission is below the detection threshold, or the knot
emission could not be reliably disentangled from the proplyd envelope.

##  3 Morphology

###  3.1 Knot identification in the MUSE data

The angular size of the giant proplyd 244-440 and the spatial sampling of MUSE
NFM+AO allow for a detailed view of the bipolar jet, which is detected in
multiple emission lines. In many cases, most notably the [O i] emission, one
observes both the jet and envelope. Lines such as [Fe ii] and [Ni ii] only
show jet emission while others (e.g., [Ar iii]) show the envelope and the
ionization front. By comparing the different morphologies traced by different
emission lines, it is possible to clearly identify the jet knots by removing
the “contamination” by nebular and envelope emission (see Figures 14$-$16).

We have identified six distinct knots with relatively strong S/N ($>10$), and
clumpy emission signatures at a low S/N ($<5$) in at least two additional
places. Not all knots are visible in our integrated channel maps due to
integration over the entire emission range, while they are more clearly seen
in smaller velocity channel maps. High velocity channels also show unresolved
emission within 0$\aas@@fstack{\prime\prime}$3 ($\sim 120$ au) of the source.
This is shown in Figures 17 and 18. We have adopted a simple notation scheme,
with each knot labeled according to its distance from the driving source, that
is to say knot E1 is furthest from the source in the eastern direction while
W3 is the knot closest to the source in the western direction. Eastern knots
are redshifted and western knots are blueshifted (see Figure 2). We list the
knot names and relative offsets from the source in Table 2.

Figure 2: Identified knots in the MUSE images. Left: The [O i]$\lambda 6300$
(top) and [Fe II]$\lambda 7155$ (bottom) emission lines with the tentative
disk and halo noted in the [O ii] panel. Emission is seen very close to the
source, but whether this is the disk or the jet is undetermined. Right: The
same images in grayscale, with the redshifted and blueshifted knots indicated
by colored circles. Table 2:  Knot positions

Knot |  $\Delta\delta$ (″) |  $\Delta\alpha$ (″)
---|---|---
E1 | $-0.5018$ | 2.7272
E2 | $-0.5338$ | 2.2182
E3 | $0.0145$ | 1.4457
W3 | $-0.0793$ | $-0.6504$
W2 |  $-$0.1006 |  $-$0.8255
W1 |  $-$0.1177 |  $-$0.9649
555Offsets of the primary knots relative to the source in 244-440 measured in
the MUSE data. The source position is taken to be 0″. Knots labeled En
correspond to eastern, redshifted knots, while Wn refers to western,
blueshifted knots.

###  3.2 Comparison with HST observations for proper motion study

Our MUSE observations were compared with archival HST data to examine the time
evolution of the outflow. Blue-shifted western jet emission is seen in the
F631N and F673N filters, and redshifted eastern jet emission is faintly seen
in the F631N filter. In both epochs of archival observations the western jet
appears as a continuous stream, making proper motion estimations unreliable
for those knots. However, the knots E1 and E3 are seen in the F631N images
with the longer exposure times.

Figure 3: Median combined F631N images of the two HST epochs used in our
proper motion study. The red contours are calculated from the residuals of the
[O i]$\lambda 6300$ Å $-$ continuum images. The “indented” structure (right
image) is due to the presence of a large diffraction spike in one of the
stacked images.

To estimate the positions of the redshifted knots, we performed centroid
fitting to the knots E1 and E3 and to the stellar profiles in both the median
[O i]$\lambda 6300$ Å and residual [O i]$\lambda 6300$ Å images using 2D
Gaussian functions. The relative offsets in the x- and y-directions are listed
in Table 3. There is the most uncertainty as to these offsets in the 1998
observations due to the low exposure times. Additionally, knot E1 is very
faint in both HST epochs, resulting in larger uncertainties also for its
position. Figure 4 shows the relative knot offsets as a function of
observation time and the best-fit line through the data, providing a proper
motion of $9.5\pm 1.1$ mas yr-1, which corresponds to $\sim 15-19$ km s-1.
This is in agreement with typical tangential velocities of HH objects in the
ONC ($<50$ km s-1; see Reiter et al., 2016) and consistent with the values
normally found in low-mass stars and substellar objects (Whelan, 2014; Riaz et
al., 2017).

Figure 4: Relative knot offsets for the three epochs of observations. A
weighted least-squares fit is plotted in red and indicates a proper motion of
$9.60\pm 0.66$ mas yr-1 for knot E3 (top) and $9.38\pm 1.54$ mas yr-1 for knot
E1 (bottom). The filled gray bars indicate the 1$\sigma$ uncertainty of the
fits.

The radial velocity can be calculated from centroid fits of the knots at each
emission line with respect to the ONC flow velocity, resulting in an average
value of $v_{\mathrm{rad}}=56\pm 10$ km s-1, and thus an absolute jet velocity
of the same order, about 60 km s-1. The resulting jet inclination angle
$i_{\mathrm{inc}}=\tan^{-1}{v_{\mathrm{rad}}/v_{\mathrm{tan}}}=$ $None$ $\pm$
$None$ with respect to the plane of the sky is not consistent with the
previous interpretations of a nearly edge-on disk, which would imply a much
lower jet inclination. We discuss our interpretation of this in Section 5.1.

These values additionally allow us to estimate a lower limit on the dynamical
age $t_{\rm dyn}$ of the knots. The knot E1 is the most distant knot with a
deprojected distance of $\sim 2\aas@@fstack{\prime\prime}8$ ($\sim$1120 au )
which yields $t_{\rm dyn}\sim 300$ yr.

Table 3:  Relative offsets and PAs

Obs. Date | E1 | E3
---|---|---
$\Delta\delta$ (″) |  $\Delta\alpha$ (″) | P.A. (°) |  $\Delta\delta$ (″) |  $\Delta\alpha$ (″) | P.A. (°)
1995 Nov 14 |  $-$0.4358 | 2.4981 | 99.9 | 0.0497 | 1.2106 | 87.7
1998 Sep 17 |  $-$0.4266 | 2.5646 | 99.4 | 0.0808 | 1.2502 | 86.3
2019 Oct 23 |  $-$0.5018 | 2.7272 | 100.4 | 0.0145 | 1.4457 | 89.4
666Offsets and PAs for knots E1 and E3 calculated from the three epochs of
observations. The offsets are measured relative to the source.

###  3.3 Jet curvature

Our data allow us to observe, for the first time, the redshifted jet and its
strong curvature. The [O i] line shows the morphology of the jet and the
direction of the ionization front. As $\theta^{2}$ Ori A is the dominant
ionization source east of the Orion Bar (O’Dell et al., 2017), one would
expect it to contribute the greatest radiation. The radiation contribution
from $\theta^{2}$ Ori B is not as powerful but is likely not negligible.
Indeed, in Figure 2 we observe knot E3 in the redshifted jet with what may be
a bow shock pointing approximately $None$ north through east, congruent with a
wind from Ori B, and the curvature traced by knots E2 and E1 appears to be in
agreement with a strong wind from $\theta^{2}$ Ori A that is deflected in part
by a lesser wind from $\theta^{2}$ Ori B.

Additionally, we see in Figure 2 a distinct asymmetry in outflow scale, with
the redshifted jet having a higher displacement than its blueshifted
counterpart. As the eastern jet flows downstream with respect to $\theta^{2}$
Ori A, it may be that crossing the photoionization front and associated
shocks, or even puncturing through the envelope, highly disturbs the western
jet.

##  4 Physical properties of the jets

The environment of the Orion region presents a unique opportunity to examine
the physical conditions of outflows and compare what we can observe in
irradiated proplyds with what is known about more isolated outflow conditions.
The electron densities and electron temperatures in the proplyd structure are
key elements for determining the mass-loss rates of proplyds, and are integral
in understanding the nature of stellar evolution (Ray & Ferreira, 2021;
Reipurth & Bally, 2001). Accurate flux measurements are thus necessary to
explore the physical conditions of the proplyd jets. In this section we
discuss the role of extinction and examine emission line ratios to explore
what diagnostic tools we can apply to this object.

In this analysis, we extracted knot fluxes using circular apertures ($r=3$
pixels or 0$\aas@@fstack{\prime\prime}$075) centered at the positions given in
Table 2. These fluxes were then corrected for extinction, as discussed in
Section 4.1. In lines where the observed flux is a combination of the knot and
envelope emission, we used apertures of the same size to extract fluxes nearby
in the envelope to estimate the proplyd contribution and subtract these from
the knots. An example of this is shown in Figure 5.

As discussed above, the physical conditions in the jet are key to
understanding the evolution of the star. The difficulty posed by proplyds,
however, is that their densities often exceed the critical limits of
traditional diagnostic ratios, such as [S ii], [O i], and [N ii] (see Bally et
al., 1998; Henney & Arthur, 1998; Mesa-Delgado et al., 2012, for a deeper
discussion). Additionally, such techniques as those proposed by Bacciotti &
Eislöffel (1999, the so-called BE99 method) rely on ratios relative to
H$\alpha$. Because of the strength of the envelope emission in H$\alpha$ in
our data, it was not possible to recover the knot fluxes in this line, and so
we could not rely on the BE99 technique.

Alternative density-sensitive lines are seen in our data, primarily forbidden
Fe lines, which can be a powerful diagnostic tool for tracing denser emission
regions (Podio et al., 2006). Additionally we observe [Ni ii] emissions, which
can be used to probe fluorescent excitation as well (Lucy, 1995; Giannini et
al., 2015).

Figure 5: Spectral profiles of the observed flux, envelope flux, and intrinsic
knot flux. The top panel compares the observed and envelope contributions,
while the bottom shows the observed and intrinsic residual. The spectra are of
knot E3 extracted from the [O i]$\lambda 6300$ Å line.

###  4.1 Extinction

Extinction plays a critical role in the analysis of the proplyd. Its
quantification is further complicated by the role of extinction in the
envelope of the proplyd as well as the photoevaporated and photoionized flow
itself (Henney & O’Dell, 1999; Mesa-Delgado et al., 2012), and the difficulty
in determining what amount of nebular emission occurs in the forefront of or
behind the proplyd envelope. Since the dust within the proplyd is not well
known, these issues make disentangling the intrinsic propyld emission from the
nebular cloud — and by extension, the jet from the envelope — a difficult
procedure.

A few methods exist to explore this phenomenon, but for more general purposes
we examined the simplest case. Using integrated channel maps of the Balmer
lines (H$\alpha$ and H$\beta$ ), we constructed maps of the extincted dust in
the FOV. This ratio was chosen due to its weak density-dependence, making it
well-suited to the analysis of the dense inner envelopes of the proplyds
(Mesa-Delgado et al., 2012). We utilized the reddening curve of Cardelli et
al. (1989) adapted for the optical regime by O’Donnell (1994) and assumed
$R_{V}=5.5$ (Weilbacher et al., 2015; McLeod et al., 2016). We further assumed
under case-B assumptions an intrinsic H$\alpha$$/$H$\beta$ flux value of 2.86
(McLeod et al., 2015) and used the observed H$\alpha$$/$H$\beta$ ratio to
produce pixel-by-pixel mappings of the $c(\mathrm{H\beta})$ reddening
coefficient along the line of sight in
PyNeb777http://research.iac.es/proyecto/PyNeb/ (see Figure 6).

What we observe is that there is a clear morphological difference between
background extinction and extinction due to dust in the proplyd envelope. In
the case of 244-440, the envelope displays a clumpy but still fairly uniform
structure, with the IF having the highest extinction. Examining the extinction
in adjacent regions and within the proplyd we generally observe a higher
extinction in the proplyd than in the nebula, which is in line with
anticipations for a dusty envelope.

With this map, we selected regions adjacent to and within the proplyd to
calculate a representative $c(\mathrm{H\beta})$ value. Under the assumption
that foreground and adjacent background extinction is uniform, we estimated
the intrinsic proplyd extinction $c(\mathrm{H}\beta)_{p}$ as the difference
between the nebular extinction ($c(\mathrm{H\beta})$=$0.44-0.45$) and that of
the proplyd envelope ($c(\mathrm{H\beta})$$\simeq 0.6$) such that
$c(\mathrm{H}\beta)_{p}=0.15$. This intrinsic reddening coefficient is similar
to what is found by Mesa-Delgado et al. (2012) for other proplyds closer to
the Trapezium.

Due to this small intrinsic extinction, we simplified our correction by taking
a constant $c(\mathrm{H\beta})$ value of 0.45 for our calculations. To
eliminate the nebular contribution, we selected a sample of background regions
away from the proplyd envelope and constructed a mean representative spectrum,
which was subtracted from every pixel in the continuum-subtracted cubes.

Figure 6: Spatial distribution map of the line-of-sight $c(\mathrm{H\beta})$
coefficient for proplyd 244-440. Figure 7: Flux ratios of the [S ii] (top),
[Fe ii] (center) and [Ni ii] (bottom) lines as a function of distance from the
source. The zero-point represents the driving source. The fluxes are summed
over circular apertures with a 3-pixel radius. Figure 8: Ratios of the [S ii]
emission line for the total observed flux (top), envelope flux (center), and
jet flux (bottom). The jet flux is estimated by subtracting nearby envelope
flux from the total observed flux and correcting for the intrinsic proplyd
extinction.

###  4.2 Shock versus photoionization

We used the fluxes measured in Table 1 to estimate line ratios and compared
them with those given in Table 2 of Giannini et al. (2015) to test for
fluorescent excitation. The predicted collisional$+$fluorescent excitation
ratio for the [Ni ii]$\lambda 7411/7378$ line is given as 0.34, while the
collisional case is given in the range 0.05$-$0.07; in our data we see ratios
from 0.13$-$0.17. The [Fe ii]$\lambda 7155/8617$ ratio is $0.8-0.9$, which is
higher than both collisional and collisional$+$fluorescence predictions.888The
$\lambda 7453$ line is omitted here due to the low S/N. These are shown in
Figure 9. These results indicate that while fluorescent pumping does appear to
enhance the emission it is not the dominant process, which is expected as
optical lines like [Fe II] are more likely to arise from collisional
excitation than photoexcitation (Bautista et al., 1996).

Figure 9: Emission line ratios of the [Fe ii] and [Ni ii] lines sensitive to
fluorescent pumping. Horizontal lines indicate the predicted ratios for each
case, taken from Giannini et al. (2015). Figure 10: Emissivity ratio of the
optically forbidden Ni emission line, calculated in PyNeb. The dashed curves
are the emissivity ratios calculated for each given temperature. The gray
region denotes the predicted range of ratios for collisional excitation. The
red region denotes the range of values we observe in our data. The solid black
line denotes the predicted ratio for collisional$+$fluorescent excitation. The
predicted values are taken from Giannini et al. (2015) assuming $T_{e}=10,000$
K and $n_{e}\sim 10^{3}-10^{5}$ cm-3.

###  4.3 Mass-loss rate in the jet

In Section 4.2 we discussed the difficulty posed by the general densities of
proplyd envelopes. Similarly, the lack of jet emission in temperature-
sensitive lines further complicates traditional diagnostic techniques,
introducing uncertainties to our determination of the electron density.
However, we note that the mass-loss rate is only weakly dependent on the
electron density and so we set reasonable bounds on this physical parameter.
We make estimates based on Figure 10, where our observed [Ni ii] ratios imply
electron densities $7\times 10^{4}$ cm${}^{-3}\leq n_{e}\leq 2\times 10^{5}$
cm-3. Our [Fe ii] ratios similarly imply electron densities $\sim$ 105 cm-3.
As the [Fe ii] and [Ni ii] ratios for the observed knots exhibit relatively
small scatter in their respective emission lines (see Figure 9), we assume a
mean density of $n_{e}=1.3\times 10^{5}$ cm-3 at $T_{e}=8200$ K.

Following Hartigan et al. (1995), we used the intrinsic knot luminosity of the
[O i]$\lambda 6300$ Å and [S ii]$\lambda 6731$ Å lines to estimate the jet
mass and outflow rate. The mass outflow (in units of $M_{\odot}$ yr-1) is
given by the equation

$\centering\begin{split}\dot{M}_{out}&=2.27\times
10^{-10}\left(1+\frac{n_{c}}{n_{e}}\right)\left(\frac{L_{6300}}{10^{-4}\leavevmode\nobreak\
L_{\odot}}\right)\\\
&\times\left(\frac{v_{\mathrm{tan}}}{150\leavevmode\nobreak\
\mathrm{km\leavevmode\nobreak\
s}^{-1}}\right)\left(\frac{\ell_{\mathrm{tan}}}{2\times
10^{15}\leavevmode\nobreak\
\mathrm{cm}}\right)^{-1}M_{\odot}\leavevmode\nobreak\
\mathrm{yr}^{-1}\end{split}\@add@centering$ (1)

for the [O i]$\lambda 6300$ Å line, where $n_{c}$ is the critical density,
$n_{e}$ the electron density, $v_{\mathrm{tan}}$ the proper motion (in km
s-1), and $\ell_{\mathrm{tan}}$ is the size of the aperture (in cm) in the
plane of the sky. The expression is similar for the [S ii]$\lambda 6731$ line,

$\centering\begin{split}\dot{M}_{out}&=3.38\times
10^{-8}\left(\frac{L_{6731}}{10^{-4}\leavevmode\nobreak\
L_{\odot}}\right)\left(\frac{v_{\mathrm{tan}}}{150\mathrm{km\leavevmode\nobreak\
s}^{-1}}\right)\\\ &\times\left(\frac{\ell_{\mathrm{tan}}}{2\times
10^{15}\leavevmode\nobreak\
\mathrm{cm}}\right)^{-1}M_{\odot}\leavevmode\nobreak\
\mathrm{yr}^{-1}\end{split}\@add@centering$ (2)

where it is assumed to be in the high-density limit so that the ratio
$n_{c}/n_{e}\ll 1$.

For these calculations, we used a tangential velocity $v_{\mathrm{tan}}=17$ km
s-1 and an aperture size
$\ell_{\mathrm{tan}}=$9.0\text{\times}{10}^{14}\text{\,}\mathrm{cm}$$. We find
low luminosities and mass-loss rates for all of the knots, with the
blueshifted jet presenting higher values due to its exposure to the ionizing
winds. For the [O i]$\lambda 6300$ Å lines the mass-loss rate is on the order
of $10^{-11}$ $M_{\odot}$ yr-1, while in the [S ii] line it is on the order of
$10^{-10}$ $M_{\odot}$ yr-1 (see Table 4).

We note however that these values do not account for photoionization effects,
and so represent a lower limit of the jet mass outflow rate. If instead we
assume photoionization plays a large role in the jets, then we can calculate
an upper limit of the mass-loss rate in the jet from

$\centering\begin{split}\dot{M}_{j}=3.4\times
10^{-9}\left(\frac{v_{j}}{100\leavevmode\nobreak\
\mathrm{km\leavevmode\nobreak\
s^{-1}}}\right)\left(\frac{n_{e}}{10^{3}\leavevmode\nobreak\
\mathrm{cm}^{-3}}\right)\\\ \times\left(\frac{r_{j}}{115\leavevmode\nobreak\
\mathrm{au}}\right)^{3}M_{\odot}\leavevmode\nobreak\
\mathrm{yr}^{-1}\end{split}\@add@centering$ (3)

as given in Bally & Reipurth (2001). With this equation, we estimate that
$\dot{M}_{j}$ would have an upper limit of $10^{-9}$ $M_{\odot}$ yr-1.

Table 4: Jet luminosity, mass, and outflow rates

Knot | [O i]$\lambda 6300$ | [S ii]$\lambda 6731$
---|---|---
$L$ | $M$ | $\dot{M}_{out}$ | $L$ | $M$ |  $\dot{M}_{out}$
E1 | $-5.2$ | $-9.0$ | $-10.5$ |  |  |
E2 | $-5.6$ | $-9.5$ | $-10.9$ |  |  |
E3 | $-5.2$ | $-9.0$ | $-10.5$ | $-5.7$ | $-9.7$ | $-9.9$
W3 | $-4.8$ | $-8.6$ | $-10.1$ | $-5.0$ | $-9.0$ | $-9.3$
W2 | $-4.7$ | $-8.6$ | $-10.0$ | $-4.9$ | $-8.9$ | $-9.2$
W1 | $-5.0$ | $-8.9$ | $-10.3$ | $-5.0$ | $-9.0$ | $-9.3$
999Luminosity, mass, and outflow rates for the redshifted knots in the [O
i]$\lambda 6300$ Å and [S ii]$\lambda 6731$ emission lines. Luminosity is
given in units of $\log_{10}$ $L_{\odot}$, masses in units of $\log_{10}$
$M_{\odot}$, and mass-loss rates in units of $\log_{10}$ $M_{\odot}$ yr-1.

##  5 Discussion

###  5.1 Implications of the proper motions

####  5.1.1 Constraints on disk inclination

Throughout our study we rely largely upon proper motion measurements to
constrain the jet velocity. This parameter is important for both
considerations of the wiggling jet model, and for calculations of the jet
mass-loss rate. In Section 3.2 we found proper motions $<20$ km s-1, and
estimated a jet inclination angle $i_{\mathrm{inc}}$ of $None$ $\pm$ $None$.
While the low proper motion is not unusual for objects in the ONC, we note
that $i_{\mathrm{inc}}$ is inconsistent with the observations of BOM2000,
which the authors interpret as a nearly edge-on disk. We examine a few
possibilities for this below.

Our assumption in the text is that the jet moves through the medium in a
ballistic manner, that is, with a constant velocity. If the jet encounters a
dense material, it may be deflected and result in a change in its bulk flow
velocity. A similar phenomenon was recorded by Hartigan et al. (2019) in the
objects HH 8 and HH 10, where stationary “loop” structures occurred in the
knots as they interacted with sheets of ambient material, resulting in no
measurable proper motions. If there are dense, unseen obstacles within the
proplyd envelope along the flow path of the jet, or if the jet is piercing
through the envelope, it may be possible that the jet is significantly slowed
or deflected as it encounters the material, or that the unshocked or weakly
shocked material in the jet produces shocks in its vicinity that falsely
present as knots. The latter case would raise the possibility that the
locations of the knots observed in the HST data are not knots at all, but
density enhancements in the surroundings. In such a scenario our proper
motions would not be reliable as a way to estimate the inclination angle of
the jet, and we would require measurements closer to the driving source that
are more likely to be unaffected by the above interactions. However, we note
in Figure 6 that while there is obvious structure in the density distribution
of the envelope, it does not appear sufficient to deflect the jet in this way.
It is important to note nonetheless that the blueshifted jet’s proximity to
the IF may affect a change in the jet structure as it passes through the
front. If this is the case, the ballistic assumption may not be appropriate
for the blueshifted jet.

In the HST images, particularly the H$\alpha$ filter, BOM2000 observed a
silhouetted structure which they interpreted as an edge-on disk (see their
Figure 7b) with a semi-major axis nearly aligned vertically in the image. In
our MUSE observations we observe a bright “halo” structure around the source
in H$\alpha$ and [O i]$\lambda 6300$ Å but we are unable to determine its
cause. If we look solely at the properties of the jet, we conclude that it
does not support an edge-on disk scenario. Moreover, if the jet is close to
the plane of the sky, the radial velocities would imply a driving source
larger than the one supported by its spectral type (see Appendix A).
Additionally, the large gap between the source and the nearest knot E3 in the
redshifted flow versus the smaller gap observed in the blueshifted side
suggests a system where the size and orientation of the disk effectively
obscures the receding jet. Finally, the jet direction is not centered with the
disk axis.

We argue that all these discrepancies are reconcilable if one assumes that the
disk observed in the HST images is not associated with the jet-driving source.
Such a scenario would explain the position angle of the jet, its low proper
motion, and our derived jet inclination angle. In our discussion on the jet
curvature below, we see that this conclusion is also a possible outcome of our
modeling and consistent with previous observations.

####  5.1.2 Dynamical age of the jet

In Section 3.2 we used the best-fit proper motion of the jet and the offset of
knot E1 from the source to present a minimum age of the jet of at least 300
yr. This is a most interesting result, as it tells us not only that the jet is
quite young and still active, but also that photoionization and
photoevaporation of the proplyd do not seem to affect the ability of the star
to launch a jet. The jet may have been launched within a “bubble” that
shielded it from the majority of hard-UV radiation, which may explain why
collisional excitation appears to dominate in the jet.

Additionally, we note above (Sec. 5.1.1) that the proximity of the blueshifted
jet to the IF may change the jet structure. Along with the age estimate, this
raises questions about the length of the blueshifted emission, which is
substantially shorter than that of its redshifted counterpart. The measured
radial velocities of the knots are not too dissimilar between the red and blue
lobes, suggesting similar launch velocities. If we further assume similar
launch epochs, then we would expect the blue lobe to extend to at least the
same length, yet we observe no emission beyond the envelope of the proplyd.
Furthermore, YSOs typically begin driving outflows early in their life-cycles
(as early as $10^{4}$ yr; see Andre et al., 2000), so it is expected that the
true extent of the jet may be far greater than what is observed (Frank et al.,
2014). This is not unreasonable, and we raise two points here to address this.

Firstly, visible knots are formed by shocks as the jet interacts with itself
or the ambient medium (in isolated regions), or by illumination from an
external source (in irradiated regions). The minimum age of the star suggests
that the jet is far older than 300 yr, so the jet must have punctured through
the envelope and passed into the larger region of the nebula, which in the
neighborhood of our object of interest typically has a density on the order of
$10^{3}$ cm-3 (Méndez-Delgado et al., 2021). Other HH objects are seen in the
ONC beyond the protective shells of the proplyds, which indicates that the ONC
is either dense enough or irradiated enough to render these objects visible.
It may be that the jet has lost so much of its density after passing beyond
the envelope that it is simply not visible, which is reasonable as the
luminosities of HH objects tend to decrease with separation from the driving
source. Secondly, the environment surrounding a proplyd is not hospitable to
transient objects. Given the low luminosity and radial velocities of the jets
we do not anticipate a very strong driving force, so it is likely that the
strong stellar winds and radiation in the ONC have entirely dissipated the
portions of the jet beyond the envelope.

###  5.2 Origin of the curvature

To explain the C-shaped symmetry common to irradiated jets in the Orion nebula
(Bally & Reipurth, 2001, and references therein), a few models have been put
forth involving ram pressure from stellar winds (Raga et al., 2009; Estalella
et al., 2012). Additionally, sinusoidal jet morphologies may arise due to the
presence of a binary companion which causes either orbital motion of the jet
source (Masciadri & Raga, 2002; Lai, 2003; Murphy et al., 2021; Erkal et al.,
2021), or a precession of the jet ejection axis due to the inner disk not
being coplanar with a companion’s orbit (Zhu, 2019; Terquem et al., 1999).
BOM2000 initially proposed the idea of a hidden companion based on the offset
of the photometric center from the geometric center of the disk as seen in
their [O i] observations. Recent evidence put forth by Tobin et al. (2009) and
Kounkel et al. (2019) also indicates that 244-440 is a spectroscopic binary.
The possibility of 244-440 possessing a companion that may produce a
“wiggling” in the jet axis, as well as the ram pressure from the stellar winds
from the stars in the Trapezium cluster and $\theta^{2}$ Ori A and B, may all
act in conjunction to produce the complex morphology that we observe in this
object.

In this section we explore whether the observed curvature in the jet can be
explained by a wiggling jet model in the absence of an appreciable side-wind.
The basic parameters of these models are the length scale of the wiggle
$\lambda$ and the half-opening angle of the jet cone $\beta$, an important
parameter in the precession model. They can be inferred by visual inspection
of the data and the proper motion estimates presented in Section 3.2. We
measured the knot centroid positions by Gaussian fitting of the jet along the
outflow axis, and utilized the jet inclination angle estimated above to de-
project these positions from the plane of the sky. In the following analysis,
we follow the method explored by Murphy et al. (2021).

From our data, we estimate a de-projected length-scale $\lambda\sim$8′′. The
half-opening angle $\beta$ is also observable from the data by fitting the
slope of the peaks of the wiggle curve. Even if we do not observe as many
peaks as some other wiggling jets, we can safely estimate $\beta=$$None$. We
can further relate the precession model to the orbital model by means of
$\beta$. In an orbital model, we define the ratio of the orbital velocity to
the jet velocity as $\kappa=v_{o}/v_{j}$, and this is related to $\beta$ by
$\kappa\leq\tan\beta$.

Using these parameters and the equations shown in Murphy et al. (2021) we
explored both an orbital motion model and a precession model for the jet. We
used the Python package lmfit101010https://lmfit.github.io/lmfit-py/ and
rewrote the equations as functions of $\lambda$ and $\beta$, allowing them to
vary by about 15%. We used a fixed value of $v_{j}=65$ km s-1 for the jet
velocity and set an upper limit on $r_{o}$ of 0$\aas@@fstack{\prime\prime}$1
(see Section 1). The results of the fits are shown in Figure 11. The errors on
the centroids were calculated according to Equation A.1 in Porter et al.
(2004). Both models were weighted with the centroid errors.

It is important to recognize that given the short length-scale of the
blueshifted jet and the inability to estimate a proper motion for that
emission lobe, it is difficult to determine whether the wiggle is point-
symmetric (precession) or mirror-symmetric (orbital motion) around the origin,
though Figure 11 does favor mirror-symmetry. Nevertheless we can make
assumptions based on the values derived from these fitted models.

A primary driver of the values derived from the fitted curves is the total
system mass. The spectral type of the visible central star implies a low-mass
object (see Appendix A) and we therefore limit our considerations to
$M_{\mathrm{sys}}$ $\leq 0.8$ $M_{\odot}$. For the orbital motion model, the
results imply $\mu$ ($=M_{\mathrm{c}}/M_{\mathrm{sys}}$) values ranging from
$0.7$ with $M_{\mathrm{sys}}$$=0.8$ $M_{\odot}$ to $0.9$ with
$M_{\mathrm{sys}}$$=0.4$ $M_{\odot}$. We can estimate the maximum binary
separation $a$ as the ratio of the orbital radius $r_{o}$ of the source about
the barycenter and the mass ratio $\mu$, $a=r_{o}/\mu$, obtaining separation
in the range $\sim 30-40$ au. The implied orbital period $\tau_{o}$ is found
by

$\centering\frac{M_{\mathrm{sys}}}{M_{\odot}}=\mu^{-3}\left(\frac{r_{o}}{\mathrm{AU}}\right)^{3}\left(\frac{\tau_{o}}{\mathrm{yr}}\right)^{-2}\@add@centering$
(4)

which suggest an orbital period of $\sim 220$ yr.

A similar calculation for the precession model results in a precession period
$\tau_{p}=238$ yr with $\mu$ ranging from $0.3$ ($M_{\mathrm{sys}}$$=0.8$
$M_{\odot}$) to $0.5$ ($M_{\mathrm{sys}}$$=0.4$ $M_{\odot}$) These results
imply orbital periods of $\sim 30-50$ yr and binary separations from $\sim
5-10$ au. This smaller value is not unreasonable as the precession model
requires a warped inner disk, and if this is induced by a companion then the
companion must be relatively close.

Figure 11: Best fits for the precession (top) and orbital (bottom) models for
the [Fe ii]$\lambda 7155$ Å jet. The relative offset along the jet axis is
deprojected by $x^{\prime}=x/\cos\theta$ for an angle of $None$ assuming a
distance $d=400$ pc. The gray region indicates a 3$\sigma$ uncertainty in the
model.

In the precession model, the inferred mass-ratio indicates an equal-mass
binary at close separation. The length scale of the deprojected jet ($\sim
4000$ au) and large fluctuation in the red lobe are consistent with precession
models, as fluctuations due to orbital motion are more likely to appear within
100s of au of the source (Masciadri & Raga, 2001). However, this model
presents some difficulties. If the companion is close in mass to the primary,
one should wonder whether or not this companion would truly be hidden in
observations. Unfortunately, we cannot resolve the binary separations in our
MUSE observations, nor is our spectral resolution high enough to perform
accurate spectro-astrometry. If we consider larger system masses, for example
$M_{\mathrm{sys}}$$\simeq 1$ $M_{\odot}$, the derived $\mu$ values still imply
an appreciably large companion. Another issue is the observed properties of
the jet. The $\mu$ values obtained with the precession model suggest a jet
driven by the primary, and this cannot be reconciled with the assumption that
the disk seen in the HST images is associated with the primary in the system.
Therefore it seems unlikely that a precession model is an appropriate
explanation for the observed wiggle.

In the orbital model, we find a mass-ratio that implies the jet-driving source
is not the primary in the system. As discussed in Section 5.1.1 this outcome
is very reasonable and is supported by observational evidence. If this is an
equal-mass binary, we would anticipate a different spectral type than what the
observed stellar spectrum suggests. If one assumes that the primary conforms
to an M0 or M1 spectral type ($M_{p}\simeq 0.5$ $M_{\odot}$), we argue that
this model gives good agreement for $M_{\mathrm{sys}}$ $\sim 0.6-0.7$
$M_{\odot}$. Additionally, the orbital model curve seen in Figure 11 best
matches the morphology of the jet. All of these arguments provide compelling
evidence that if the curvature can be explained by a wiggling jet model, the
orbital motion model is a strong candidate, and that the jet is associated
with the smaller, hidden star in the system.

We have not explored the impact of stellar winds here. Generally, side-wind
deflection models imply a hyperbolic curvature as the jet is deflected away
from the wind source (Raga et al., 2009), although more complicated
morphologies in photoionized regions are possible (Masciadri & Raga, 2001). In
our observations, the C-shaped morphology anticipated by a deflection model
appears to be applicable primarily to the envelope but not the jet, indicating
that the jet is either shielded to some degree from the winds or that multiple
winds are influencing the system in a way that is beyond the scope of this
paper. Furthermore, using the derived proper motions and overlaying their
vectors on the data (see Figure 12), it is seen that the knots do point
radially outward from the source with a slight difference ($\sim$$$) in their
directions. The decrease in angle exhibited by knot E1 could potentially
indicate an interaction between the jet and a side-wind; however with proper
motions available for only two knots either conclusion is approached with
caution. Nevertheless these models lay a positive groundwork, as future
observations may help provide greater constraints on the parameters of a
potential companion.

Figure 12: Proper motion vectors superimposed on the [O i]$\lambda 6300$ Å
emission line. A small difference is observed, pointing overall radially
outward though possibly pointing along the direction of the jet curvature. The
angles are measured east of north in the plane of the sky.

###  5.3 Diagnostics and mass-loss

Even if we could not reliably determine the temperature of the jet in our data
through traditional diagnostic diagrams, the presence of refractory species
like [Fe ii] and [Ni ii] provides opportunity to explore ranges of density, as
well as the potential role of fluorescent pumping of the jet. It is reasonable
to assume that the presence of external ionizing sources may induce
fluorescence in the jet, particularly given the presence of the [Ni ii] lines
and the observed line ratios in Figure 9. However, the density ranges
estimated from the [Ni ii] lines ($\sim 10^{4}-10^{5}$ cm-3) are above the
critical limit for fluorescent pumping, which suggests that shock processes
play the larger role in producing the emission lines. This may be due to the
envelope shielding the jet to some degree from incident radiation, making it
not as exposed as others to the external environment.

This result has interesting implications when considering the mass-loss rate.
In Section 4.3 we presented limits on $\dot{M}_{j}$ in the range
$10^{-11}-10^{-10}$ $M_{\odot}$ yr-1 and argued that even if photoionization
plays a large role, we would not expect a mass-loss greater than $10^{-9}$
$M_{\odot}$ yr-1. Additionally there is an asymmetry in the luminosities and
mass-loss rates with the blueshifted jet presenting generally higher values.
While the jet may benefit from shielding by the envelope, the fact that the
computed values are higher in the blue lobe than the red demonstrates that the
jet may still be partially exposed to external radiation. These lower values
are typically seen in very low-mass protostars and brown dwarfs (Whelan et
al., 2009; Riaz et al., 2017), which is consistent with the spectral typing of
the source (Figure 13).

##  6 Conclusions

We have performed the first analysis of the possible origins of the curvature
and physical conditions of the proplyd 244-440 jet using high spatial
resolution IFU observations with MUSE NFM+AO. We identifed in the MUSE data
multiple, previously unreported knots in the redshifted jet lobe. These were
observed in various emission lines, most notably [O i] and refractory species
like [Fe ii] and [Ni ii]. Two knots (E1, E3) are also visible in archival HST
images.

Using measurements of the E1 and E3 knots in the MUSE and HST data, we
estimated a low proper motion of $9.5\pm 1.1$ mas yr-1 and inferred a jet
inclination angle $i_{\rm jet}=$$\pm$$$, which appears contrary to previous
interpretations of a nearly edge-on disk. Closer analysis suggested that the
jet is not associated with the observed disk, and we posit that the jet is
actually driven by a smaller companion.

We utilized a jet-wiggling model to explore for the first time the curvature
in the jet, and found that in the absence of other forces (i.e., multiple
strong side-winds) the curvature could be explained by orbital motion of the
jet source. As recent evidence suggests this is a spectroscopic binary, this
is not unreasonable. We further reason that if this is due to a companion, we
might expect the driving source to be $\leq 0.15$ $M_{\odot}$ in orbit around
an M0 or M1 type star ($M\sim 0.5$ $M_{\odot}$) at a separation of $\sim 30$
au.

Using the [O i] and [S ii] lines, we estimated a lower limit on the mass-loss
rate in the jet on the order of $10^{-10}-10^{-11}$$M_{\odot}$ yr-1. If we
assume the jet is nearly completely photoionized we set an upper limit on the
mass-loss rate of $<10^{-9}$ $M_{\odot}$ yr-1. We note that similarly small
values are observed in low-mass and substellar objects such as brown dwarfs.

Finally, we looked at the “proplyd lifetime problem” and estimated an
evaporation time between $0.1-0.2$ Myr. The minimum dynamical age of the jet
was found to be around 300 yr, indicating that the source is still quite
active and that the jet may still be quite young. Compared to the evaporation
time of this proplyd, this also tells us that photoionization and
photoevaporation of the proplyd had likely been occurring for some time before
the jet was launch. This might indicate that the envelope has acted as a
protective shell enclosing the jet and shielding it from a significant portion
of Ly$\alpha$ radiation. This yields critical information about the durability
of the accretion-outflow connection in the harshest of conditions. We also
raise the possibility that the calculated dynamical age may drastically
underestimate the true age, reasoning that if the jet has extended beyond the
envelope it may have been completely destroyed.

This work demonstrates the power of the VLT/MUSE NFM instrument in exploring
jet launching dynamics. The data yielded by this instrument is rich in
information, and is capable of exploring spatial structure across many
emission lines critical in the study of stellar jets. The high angular
resolution provided by the instrument is particularly ideal for the
exploration of jets in high-radiation environments, as the external
irradiation exposes quiescent, unshocked material, and allows us to better
identify emission features and remove nebular and envelope contributions to
the jet emission.

###### Acknowledgements.

We are grateful to the referee Alex Raga for his feedback which has helped us
improve the quality of this work. A.K. would like to acknowledge funding
through the John and Pat Hume Doctoral Scholarship at Maynooth University,
Ireland. C.F.M. is funded by the European Union under the European Union’s
Horizon Europe Research & Innovation Programme 101039452 (WANDA). S.F. is
funded by the European Union under the European Union’s Horizon Europe
Research & Innovation Programme 101076613 (UNVEIL). Views and opinions
expressed are however those of the author(s) only and do not necessarily
reflect those of the European Union or the European Research Council. Neither
the European Union nor the granting authority can be held responsible for
them. We also extend our thanks to Monika Petr-Gotzens and Teresa Giannini for
their helpful comments and discussion.

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##  Appendix A Supplemental Information

Many of the results discussed in this paper require us to have at least a
broad estimate of the source mass. If we are to utilize a jet wiggling model,
for example, a key assumption is that the source contains a binary as a
critical parameters is the ratio of the companion mass to the primary.
Similarly, any jet proper motion study or computation of mass-outflow rates
must be compared against some mass if we are to determine how reasonable our
values are. To accomplish this, we extracted an on-source spectrum, corrected
for extinction, and compared it with several spectra of known stellar types to
find which is most similar as shown in Figure 13.

Our spectral templates were all observed on the X-Shooter instrument and their
spectral types reported in Manara et al. (2013) and Manara et al. (2017).
Based on this, we argue that 244-440 is most likely an M0 or M1 spectral type
star, placing it on the very low-mass end. We believe this to be reasonable as
well, as the proper motions and radial velocities are both quite low,
indicative of a low-power outflow. Additionally, the estimated mass-outflow
rate is comparable with those seen in brown dwarfs, strengthening this
argument (Whelan et al. 2014; Riaz & Whelan 2015; Riaz et al. 2017).

Figures 14$-$16 show flux-integrated three-color composites for various
emission lines extracted from both background-subtracted and unsubtracted
cubes. Figures 17 and 18 show velocity channel maps isolate the redshifted and
blueshifted emission features. Figures 19 and 20 show these features in color
composites using flux-integrated spectro-images.

Figure 13: On-source spectrum of 244-440 (red) overplotted with spectra from a
sample of M-type YSOs. The spectral type of each YSO is given in the legend of
each panel.

Figure 14: Three-color flux-integrated image composites of the proplyd 244-440
for the background-subtracted (left) and unsubtracted (right) data. Red is [Ar
III]$\lambda 7136$, green is H$\alpha$, and blue is [O i]$\lambda 6300$.

Figure 15: Same as Figure 14. Red is [Fe ii]$\lambda 7155$, green is [N
ii]$\lambda 6548$, and blue is [O i]$\lambda 6300$.

Figure 16: Same as Figure 14. Red is [O ii]$\lambda 7320$, green is [O
i]$\lambda 6300$, and blue is [O iii]$\lambda 5007$.

Figure 17: Velocity channel maps of the [O i]$\lambda 6300$ line in the MUSE
data.

Figure 18: Same as Figure 17, but for the [Fe ii]$\lambda 7155$ line. Figure
19: Three-color composite of the [O i]$\lambda 6300$ emission line. Green is
the full flux-integrated image, and red and blue are the red- and blueshifted
emissions shown in Figure 17. Figure 20: Three-color composite of the [Fe
ii]$\lambda 7155$ emission line. Green is the full flux-integrated image, and
red and blue are the red- and blueshifted emissions shown in Figure 18.