Imaging of HH80-81 Jet in the Near-infrared Shock Tracers H₂ and [Fe ii]

We imaged the IRAS 18162-2048 region using the Wide-Field Camera (WFCAM; Casali et al. 2007) mounted on the 3.8 m United Kingdom Infrared Telescope (UKIRT). The WFCAM has four 2048 × 2048 HgCdTe Hawaii-II arrays with an optical system that provides a pixel scale of 04. This provides a total field of view of 135 × 135 per array and 0.21 deg² for all four arrays together. The target was located in the center of one of the four arrays and the data from only that array were used in the current study. The observations were performed by dithering the target to nine positions on the array and a 2 × 2 microstepping was done. The resulting mosaics have an image scale of 02 pix⁻¹. Observations were carried out using the NIR broadband J, H, K filters and the narrowband H₂ and [Fe ii] filters. The narrowband H₂ filter is centered at λ = 2.122 μm and the [Fe ii] line filter at λ = 1.644 μm with a FWHM of 0.021 μm and 0.028 μm, respectively. The observations were conducted between 2020 February 20 and 2021 September 20.

The preliminary reduction including the creation of mosaics was carried out by the Cambridge Astronomical Survey Unit (CASU). For the narrowband H₂ and [Fe ii] filters, further reduction including continuum subtraction was carried out using the Starlink software (Currie et al. 2014). We have obtained imaging observations of H₂ in four epochs and those of [Fe ii] in six epochs. The total integration time for J, H, K, H₂, and [Fe ii] filters were 720 s, 360 s, 180 s, 1440 s, and 1440 s, respectively, for each epoch of observation. For each epoch, the H₂ and [Fe ii] images were aligned with K- and H-band images, respectively, using WCSALIGN. Due to variations in the seeing conditions between the narrowband and broadband observations, images with smaller point-spread functions (PSFs) were smoothed to match with the FWHM of the images with larger PSFs using the task GAUSMOOTH. For continuum subtraction, a sufficient number of isolated bright point sources were selected and their flux counts from both the narrowband and broadband sky-subtracted images were calculated. The broadband to narrowband flux ratio was then computed. For each imaging observation in an epoch, the broadband image was scaled using the average value of the ratio of these fluxes and subtracted from the narrowband image to obtain the continuum-subtracted image for that epoch. The continuum-subtracted images of all the epochs were averaged to obtain the final image.

The flux calibration of the continuum-subtracted images was carried out using the H- and K-band flux densities of isolated point sources in the field. For this, we first interpolated the broadband H- and K-filter flux densities of a few isolated point sources to the central wavelength of the narrowband filters. These flux densities were multiplied by the FWHM of the narrowband filters to obtain the total fluxes. We then derived the photometry of the point sources on the normalized narrowband images to obtain the counts in ADU. This was used to estimate the flux per ADU, and this was used to scale the counts in ADU obtained from the continuum-subtracted narrowband images to estimate the fluxes. The fluxes were not corrected for extinction, as the extinction value toward each knot is not known.

In addition to this, the (3.678 μm) and (4.686 μm) imaging observations were carried out using the UKIRT 1–5 μm Imager Spectrometer (UIST; Ramsay Howat et al. 2004). UIST has a 1024 × 1024 InSb array. The observations were carried out using the 012 pixel⁻¹ camera. The object was dithered to four positions along the corners of a square of 40″ side. The total integration time for the and observations were 160 s and 280 s, respectively. For each setting, a dark observation was done, which was followed by the science observations. The dark-subtracted science frames were median combined to create sky flats. Frames in individual pairs were subtracted from each other and divided by the flats. This resulted in a positive and negative beam in the mutually subtracted, flat-fielded image. The final mosaic was constructed by combining the positive and negative beams.

The optical emission from the HH80-81 jet in the filters Hα + [N ii], [S ii], and [O iii] has been imaged using the Wide Field and Planetary Camera 2 (WFPC2) on board the Hubble Space Telescope (HST). These high angular resolution images have been published by Heathcote et al. (1998). We have extracted these images from the Hubble Legacy Archive. We have employed these images to compare the morphologies of HH80 and HH81 with those from our NIR observations, in order to characterize the nature of shocks.

The full length of HH80-81 jet as observed in H₂ emission is shown in the central panel of Figure 1. The surrounding panels display the emission toward the jet in H₂, [Fe ii], and the optical Hα + [N ii] filters in red, green, and blue, respectively, revealing the morphologies of individual knots. The details of the sizes and locations of the subimages are listed in Table 1. We observe that in the image obtained by us, the jet displays a chain of knots extending up to 1.9 pc in the northeast and 2 pc in the southwest directions, as projected on the sky plane. Considering the inclination angle of the jet to be 56° (Heathcote et al. 1998) from the sky plane, the true extent of the bipolar jet as imaged in NIR is ∼7 pc. It is possible that there are additional knots extending farther away from this toward HH80N. We identify a total of nine new knots along the HH80-81 main jet direction and multiple arcs in H₂ emission. Four knots are in the northern arm and five are in the southern arm of the jet. We assign numbers 1–9 to the knots as seen in our NIR images from north to south. The Knots 8 and 9 correspond to HH81 and HH80, respectively. We describe a few morphological features of the knots below:

• 1.  
  Knots 1 and 2 exhibit only H₂ emission, and Knot 4 is revealed in only [Fe ii] emission, while the remaining knots exhibit emission in both [Fe ii] and H₂ transitions.
• 2.  
  Knots 1–6 are elongated along the direction of the jet. Knots 7–9 appear broken and have complex morphologies.
• 3.  
  In none of the knots, the regions corresponding to [Fe ii] and H₂ emission completely overlap. In Knots 3 and 6, the [Fe ii] emission clearly lies ahead of the H₂ emission along the outward jet direction; whereas, in the case of Knot 5, the reverse is seen.
• 4.  
  The NIR line emission for Knots 8 and 9 do not coincide with optical line emission. For other knots, the optical data are not available.

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The central region close to the driving source of the jet is shown in Figure 2. The Spitzer 8 μm broadband image is shown in red and the H₂ emission is depicted in green. The maps are overlaid with the contours of [Fe ii] emission in orange. The association with radio emission is illustrated by contours (cyan) corresponding to the 610 MHz emission (Vig et al. 2018). We observe that the emission in H₂ is widely distributed across the image, unlike the [Fe ii] emission. It is also evident that the [Fe ii] emission arises from the highly ionized inner jet close to the axis, while the mid-infrared (MIR) emission delineates the outflow cavity. Multiple bow shocks and knots are observed in H₂ and the majority of this emission appears to be distributed in directions away from the HH80-81 jet axis. The radio jet is well delineated by the [Fe ii] emission.

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The prominent H₂ features in the region are:

• 1.  
  The HH80-81 main jet: This extends along the northeast–southwest direction. The emission is seen as knots extending along the jet as well as arcs close to the central region. A detailed discussion of this jet is presented in Section 4. The catalog number of the knots associated with this jet in the Catalogue of Molecular Hydrogen Emission-Line Objects (MHOs) in Outflows from Young Stars (Davis et al. 2010) is MHO 2354.
• 2.  
  Reg 1: This comprises a chain of five H₂ knots toward the eastern side, which are enclosed within a dashed ellipse shown in Figure 2.
• 3.  
  Reg 2: This includes two distinct H₂ knots to the northeast of Reg 1.
• 4.  
  Reg 3: This region includes arcs of H₂ emission along the east and west directions on either side of a YSO (Qiu et al. 2008), as shown in Figure 2, facing inwards of the dashed ellipse of size . The arc toward the east is very strong in emission and with overlapping emission at 8 μm, it is seen as yellow in the figure. We note that the nebulosity at 8 μm appears well enclosed within the defined ellipse.
• 5.  
  Reg 4: This region shows two knots aligned northeast with respect to the central region.
• 6.  
  Reg 5: The H₂ emission shows arc-shaped patterns over a region of size . The arcs appear to surround the nebulosity observed at 8 μm. The emission to the south of the enclosed ellipse shows strong H₂ emission, overlapping with 8 μm emission.
• 7.  
  Reg 6: This region includes a knot toward the west of the central region.
• 8.  
  Reg 7: Two faint arcs in H₂ are observed facing each other within the region enclosed by the ellipse in this region.

The MHO catalog numbers of Regs 1–7 are MHO 2355−2361. A detailed description of the morphology of knots in narrow-line emission bands is presented in Appendix A.

We have estimated the integrated fluxes of knots in the narrowband H₂ and [Fe ii] filters by using elliptical apertures for each of the knots, which are marked in cyan in Figure 1. The knot fluxes along with the sizes of apertures (semimajor × semiminor) used for the flux calculation are tabulated in Table 2. The H₂ and [Fe ii] fluxes of the knots are in the range (0.4–5.2) × 10⁻¹⁴ erg s⁻¹ cm⁻² and (3.1–13.6) × 10⁻¹⁴ erg s⁻¹ cm⁻², respectively. Knots 1 and 2 do not have any associated [Fe ii] emission. We have also calculated the ratio of [Fe ii]/H₂ emission toward knots where both are detected. This ratio is represented as , and estimated in order to understand the nature of shock across different knots. These values are also tabulated in Table 2 and are comparable with those obtained from observational studies conducted toward other YSO jet sources (Lorenzetti et al. 2002; McCoey et al. 2004; Garcia Lopez et al. 2010). It is worthwhile to note that the extinction can be quite nonuniform along the length of the jet, and will also be affected by the inclination of the jet with respect to the line of sight. However, having stated this, we observe that varies widely across the knots, from 0.6 in Knot 7, suggesting relatively low excitation of [Fe ii] emission, to 12.1 in Knot 5, where [Fe ii] emission clearly dominates over the H₂ emission. This clearly indicates that there is a considerable difference in the excitation conditions of the knots. We discuss this in detail in Section 4.

Table 2. Fluxes of HH80-81 Jet Knots, the Sizes of Apertures for the Flux Calculation, and the [Fe ii]/H₂ Line Ratios ()

Source Name	H2 Flux	[Fe ii] Flux	Aperture Size	= [Fe ii]/H2
	(10−14 erg s−1 cm−2)	(10−14 erg s−1 cm−2)
Knot 1	1.7	⋯ a	107 × 47	⋯
Knot 2	1.3	⋯ a	81 × 47	⋯
Knot 3	2.6	4.6	148 × 55	1.8
Knot 4	⋯ b	11.7	139 × 60	⋯
Knot 5	0.4	4.6	107 × 39	12.1
Knot 6	0.9	5.2	59 × 58	6.0
Knot 7	5.2	3.1	146 × 107	0.6
Knot 8	1.5	7.0	96 × 91	4.7
Knot 9	4.3	13.6	162 × 121	3.2

Notes.

^a [Fe ii] emission is not detected in these knots. ^b H₂ emission is not detected in this knot.

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In this section, we describe the emission from the central region in the NIR and MIR broadband filters. The emission through the NIR J, H, and K bands can be seen in Figure 3. The J-, H-, and K-band emission toward the northern extension of the jet is up to pc. That said, toward the southern arm, the NIR J- and H-band emission is observed up to 0.2 pc, whereas the K-band emission is seen to extend up to 0.3 pc. We identify a sharp bend in the jet cavity toward the northern arm, where it is initially directed westwards but veers toward the east; see Figure 3. The most massive millimeter cores observed in this region are MM1 and MM2 (Busquet et al. 2019), which are marked in the figure. In addition, the knots and arcs that are discussed in Section 3.1 are also shown in the figure.

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The emission toward the central region in the and bands can be viewed in Figure 4. We note that the cavity is not seen very clearly in these images and we attribute this to the low signal-to-noise ratio images. While the -band image shows hints of emission toward the jet cavity near the center (see Figure 4(a)), this is not detected in the image. A comparison with MIR IRAC images from the space-based Spitzer mission shows the presence of strong MIR emission from the cavity walls (Figure 2). Cavities around jets in the MIR have been observed earlier and the emission has been attributed to thermal dust, which is directly heated by the central source that drives the outflow (De Buizer 2006). We observe a few point-like sources in our moderately high resolution and images, and we identify a few of them as cores that have been observed using millimeter ALMA observations (Busquet et al. 2019). The brightest emission in the image is observed at ∼2000 au to the northeast of MM1. While this emission is very likely from the outflow cavity, there exists a possibility that this is a YSO, considering the point-like morphology seen in the image. The lack of associated millimeter emission gives more credence to the outflow cavity hypothesis, although the possibility of an evolved YSO cannot be completely ruled out. MM1 is brighter in the band compared to band, indicating that it is indeed young. We have not detected any emission associated with MM2(E) and MM2(W) in the band.

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In this section, we qualitatively characterize the nature of shocks at play in each of the knots. Depending on the strength, structure, and nature of the shock generated by jets, the shocked plasma can cool via emission of atomic, molecular, or ionic lines. Thus the emission-line features can be utilized to uncover the type and nature of shocks that give rise to these emissions. Based on their physical properties, shocks can be broadly classified as (i) Jump shocks (J-type), (ii) Continuous shocks (C-type), and (iii) J-type with a magnetic precursor (Draine 1980; Flower et al. 2003). C-shocks are typically associated with velocities of ∼40–50 km s⁻¹, low ionization of molecular gas, and strong magnetic fields. They are rich in infrared emission lines (H₂, CO, H₂O, etc.). J-shocks, by contrast, are faster and mostly rich in optical and ultraviolet emission lines and require high ionization levels with low magnetic field strength (Hollenbach & McKee 1989). Younger shocks typically tend to be of J-type (strong), and such shocks may eventually evolve into C-type (weak) in the presence of a strong magnetic field. A shock that has not attained thermal equilibrium could display the properties of both C- and J-type shocks. The intermediate stage in this evolution is called the J-type shock with a magnetic precursor. While spectroscopy is ideal for probing the intensities of excitation lines and hence the nature of shock at a specific location, the spatial distribution of the different excitation lines through a morphological study is equally important, as it provides strong clues about jet propagation in the ambient medium.

We utilize the H₂ and [Fe ii] line images to comment on the possible nature of shocks in different knots of the HH80-81 jet. Molecular H₂ is a good tracer of low-velocity (10–50 km s⁻¹) weak molecular shocks associated with protostellar jets, i.e., shocks at low velocities (Draine 1980). In such cases, the dissociation of the H₂ molecule does not occur. However, in regions of low densities, the H₂ molecule remains undissociated even up to 80 km s⁻¹ (Le Bourlot et al. 2002). Infrared H₂ emission could occur either in nondissociative C-shocks or low-velocity J-shocks. This is corroborated by the fact that numerous HH flows (Schwartz et al. 1988; Lane 1989) and several CO outflows have exhibited H₂ emission features tracing weaker regions of shocks with low excitation levels. J-shocks with speeds greater than 25 km s⁻¹ are found to dissociate H₂ (Kwan et al. 1977), whereas C-shocks are capable of accelerating H₂ to velocities of up to 50–80 km s⁻¹ without dissociating it (Draine 1980; Le Bourlot et al. 2002). [Fe ii] emission is widely used to understand fast and dissociative J-shocks caused by jets with velocities larger than 50 km s⁻¹. At these velocities, the head of the bow shock can destroy the grain in which iron is locked up (Lorenzetti et al. 2002). The grain destruction releases iron in the gas phase and this is followed by ionization of the neutral iron through charge transfer reactions with ions. It is possible that J-shocks can have cooler wings, where H₂ molecules are excited behind the [Fe ii] emitting region. Alternately, J-shocks with magnetic precursors can give rise to a geometry in which H₂ emission is seen in front of the [Fe ii] emission (Draine 1980). The morphology of emission in shock tracers can provide a resourceful gauge to examine the interaction of the jet with the ambient medium.

In the HH80-81 jet, from the observed morphology of knots aligned along the jet, we attempt to segregate strong and weak shocks and correlate this information with signatures of C-shocks and J-shocks, as well J-shocks with magnetic precursors. In Knot 1, the elongated cavity-like H₂ emission and the absence of [Fe ii] and radio emission toward this distant knot could be explicable by the weakening of the shock strength with radial distance as it moves outwards from the central source. The absence of [Fe ii] and radio emission helps to exclude a J-shock as the exciting mechanism at this knot location, and hence we propose that the shock in this knot could be a C-type shock. For Knot 2, although its morphology is different from that of Knot 1, strong H₂ emission and the absence of [Fe ii] and radio emission again points toward the possibility of a C-type shock due to weakening of the shock with radial distance from the central source. The filled morphology could suggest the presence of a stronger shock. Compared to Knot 1, this shock is possibly stronger due to the smaller radial distance from the driving source. Alternatively, the emission could be from the front side of the wall cavity as mentioned earlier.

In Knot 3 the [Fe ii] emission appears to arise from the head of the bow shock and the H₂ emission corresponds to the cooler wing of the bow shock, i.e., interior to that of [Fe ii] emission in the postshock cooling zone. The fast dissociative shock head of the J-shock results in the generation of Fe⁺ ions in the immediate postshock region, which undergoes collisional excitation and emission of the forbidden [Fe ii] transition. Behind the shock front, molecular reformation occurs due to which H₂ emission is observed (Hollenbach & McKee 1989). For the H₂ molecule to remain undissociated, the angle between the direction of the oblique planar shock responsible for the emission and the jet direction is expected to be less than 5°–7° for a J-type shock and 10°–15° for a C-type shock (Davis et al. 2000). In Knot 3, we find that the angle between the H₂ knot and the flow direction is ∼16° and between the [Fe ii] knot and jet is ∼42°. This strongly suggests the C-type nature of the cooler wings and J-type nature of the shock head where H₂ emission is missing. A similar case is observed toward Knot 6. The [Fe ii] emission is located at the bow head and H₂ emission is at the cooler wings behind the shock front. This again points to the possibility of a strong dissociative J-shock with cooler wings. In Knot 4, only [Fe ii] emission is present due to the highly dissociative nature of the shock.

We observe the presence of H₂ emission exterior to the [Fe ii] emission in Knot 5, and it is possible that a J-shock with a magnetic precursor can account for this scenario. The magnetic precursor causes heating and compression of the materials upstream of the shock front when the shock wave travels in a weakly ionized gas in the presence of a transverse magnetic field (Hartigan et al. 1989; Flower et al. 2003). However, here the neutrals undergo a discontinuous change of physical conditions across the shock. A low ionization fraction reduces the rate of ion–neutral collisions, resulting in weaker compression of the magnetic field. Under these low ionization conditions, if the magnetic field is strong enough the shock will possess a magnetic precursor, which compresses the magnetic field even before the shock front crosses a region. This happens in this scenario because the magnetosonic waves can propagate faster than the actual shock front (Draine 1980). Hence, the arrival of shock compression before the shock front at each location along the path of the jet flow provides favorable conditions for excitation of H₂ molecules without dissociating it. This could explain the observed H₂ emission in this knot. Furthermore, behind the shock compression, at the dissociating shock front, we observe [Fe ii] emission in accordance with predictions from the theory.

The Herbig–Haro objects HH80 and HH81 display complex morphologies. For HH81 (Knot 8), the H₂ and [Fe ii]/optical emission appear to arise from bow shocks that are arched in a direction that is opposite to the jet propagation. The arching of the bow shock opposite to the direction of the flow can be explained by a scenario where the jet encounters a small and dense clump of material on its path (Stapelfeldt et al. 1991). The bow shock then curves in the direction opposite to the direction of the jet flow. Toward HH80 (Knot 9), multiple knot-like structures are observed. The presence of elongated H₂ emission toward the eastern lateral edge of the [Fe ii] emitting region could have an origin in the wiggling or sideways motion of the jet flow (Marti et al. 1993; Eisloffel et al. 1994). An alternative explanation is provided by Noriega-Crespo et al. (1996), who introduce the grazing jet/cloud core collision scenario. This states that when the jet impacts on the molecular core at an angle, it gets deflected and moves away from the surface of impact. As the jet deflects away, the external pressure forces it back, causing the outflow to bounce back and forth while some gas slides closer to the core. This is the zone of interaction between the atomic and molecular gas, and the H₂ excitation most likely arises from this region. We suggest that the back-and-forth motion of the jet could create internal shocks that are substantiated by the presence of elongated [Fe ii] emission, located closer to the jet axis than the H₂ emission. The complex morphology of the entire knot could arise from a scenario where thermal instabilities develop in an optically thin jet that ultimately fragments and breaks up the shocked material (Field 1965; Hunter 1970). Proper-motion studies aimed at understanding the kinematics of these HH80-81 knots have measured velocities for all the subknots, and these lie in the range 74–916 km s⁻¹ (Heathcote et al. 1998). The direction of velocity is different for different subknots, which supports our claim regarding fragmentation of these knots as a result of thermal instabilities. For HH80 and HH81, it is difficult to conclude the nature of shock that explains the entire complex morphology of knots using the NIR and optical alone, and it is possible that one or multiple types of shocks could be generated in the subknots.

The value of each knot is representative of the strength of the shock generated. The Hollenbach & McKee (1989) model predicts that a [Fe ii]/H₂ ratio of greater than 0.5 is indicative of high-velocity dissociative shocks. The values of the HH80-81 jet knots are greater than 0.5 (Table 2), thereby indicating the presence of J-type shocks throughout the length of the jet. In the case of Knot 5, which is located in the southern arm of the jet, the value is significantly high, due to its close proximity to the driving source. Here, the shocks are very strong and highly dissociative, thereby favoring [Fe ii] emission over H₂. Similarly, the corresponding closest knot in the northern arm, Knot 4, is much stronger and dissociative, resulting in only [Fe ii] emission.

Next, we utilize the location of soft X-ray emission to discuss the plausible nature of shocks in these regions. The highly dissociative head of a J-shock, being the strongest section of the shock, is believed to be rich in UV and X-ray emissions. X-ray emission has been detected toward HH80 (subknots A, C, and G) and HH81 (subknot A) by Rodríguez-Kamenetzky et al. (2019) that has overlap with the optical emission. In addition, a scrutiny of the [Fe ii]/optical and soft X-ray images (Figure 2 in Rodríguez-Kamenetzky et al. 2019) suggests the X-ray emission to be marginally ahead of the [Fe ii]/optical emission in the outward propagating jet direction. This points toward the likely presence of a J-type shock at these locations. For the other subknots, the shock mechanisms are difficult to disentangle.

Although we have tried to understand the nature of shocks, i.e., strong (J), weak (C), or intermediate (J with magnetic precursor), based on the spatial distribution and overlap of H₂ and [Fe ii] lines, it is possible that extinction effects could also play a role in regions where H₂ is detected but [Fe ii] is not detected, as the former suffers lower (∼50%) extinction compared to the latter. This is discussed in detail in Section 4.2.

A comprehensive view of the jet reveals that toward the northern arm, the knots closer to the driving source display solely [Fe ii] and radio emission or a combination of [Fe ii], radio, and molecular H₂ emission, whereas H₂ emission dominates in the farther knots. The H₂ emission is believed to trace the diffuse regions in outflow lobes that are not well collimated, whereas [Fe ii] and radio emission reveal the compact jet where the fast moving jet ejecta produces stronger shocks (Lorenzetti et al. 2002). This implies that [Fe ii] and radio emission trace the ejecta accelerated directly by the exciting source, whereas H₂ emission arises from shocked ambient gas where the ejecta interacts with it. The [Fe ii] emission gradually diminishes with radial distance from the exciting source toward this arm. From this, we can infer that at smaller radial distances from the driving source, shocks are stronger J-type, and they appear to gradually weaken farther along the jet length.

In the southern arm, all the knots show both [Fe ii] and H₂ emission. HH81 and HH80, located toward the southern extremity of the jet, are complex in nature, which we attribute to thermal instabilities associated with low densities. Overall, the nature of shocks in the knots would suggest that densities of the ambient medium are relatively lower toward the southern arm as compared to the northern arm. This is validated by the fact that optical emission from HH81 and HH80 is visible due to lower extinction. This is also in accordance with conclusions of Marti et al. (1993) that the knots in the northern arm interact with the molecular cloud, unlike the southern knots HH80 and HH81, which lie beyond the edge of the molecular cloud. A comparison of fields toward the northern and southern regions of the jet in the JHK three-color-composite image shown in Figure 5 of Appendix B indicates that the northern field displays a lower number of sources that are redder compared to those in the southern field.

In view of the fact that we detect both [Fe ii] and H₂ emissions in all the knots in the southern arm, we carefully analyze the variation of the corresponding values with radial distance. The values for all the knots are tabulated in Table 2. One can discern that, with an increase in radial distance from the central YSO, decreases from Knot 5 to Knot 9 with a sharp decline in the value at Knot 7. The decrease in with radial distance is likely to be due to a decrease in the strength of dissociative shocks. Such a decrease in with radial distance has also been observed in other jet systems (Davis et al. 2000; Lorenzetti et al. 2002; Nisini et al. 2002). In addition, the lower densities toward Knots 8 and 9 could also result in a decrease in H₂ emission in these knots compared to the ones at smaller radial distances.

For each knot, we compare the NIR narrow-line observations with the radio emission from ionized gas, wherever available. We note that the knots giving rise to ionized gas emission are associated with [Fe ii] emission. However, the converse is not true, and the exceptions are Knot 5 (no 610 MHz emission) and Knot 7 where the [Fe ii] emission does not have a corresponding ionized gas emission. This is consistent with our assertion that the shock in this knot is J-type with a magnetic precursor. This is weaker than the strong J-type shock and it is highly likely that ionization of Fe is favorable against the ionization of hydrogen, as the former has a lower ionization potential than the latter. This would explain the lack of ionized gas emission as well as the presence of [Fe ii] toward this knot.

The overall geometry of the jet on a large scale shows a curved S-shape centered on the MM1 core in IRAS 18162-2048. This S-shape of the distribution of knots was first noted in radio by Marti et al. (1993). We have used NIR emission from multiple knots including those that emit in radio to confirm the curved morphology of the knot train. In addition, we note that a direct evidence of wiggling is seen in HH80 (Marti et al. 1993), where H₂ is linearly elongated along the lateral edge of the jet, which provides support for a wiggling jet scenario. Similar geometries have been observed before in various other jets (Eisloffel & Mundt 1997; Gomez et al. 1997; Lorenzetti et al. 2002). Two possible scenarios can be considered: (i) the jet interacting with dense ambient medium (Lorenzetti et al. 2002), or (ii) the wiggling of the jet (Raga et al. 1993). The symmetric shape is believed to result from the wiggling of the jet due to precession (Bachiller et al. 2001; Lorenzetti et al. 2002; Takami et al. 2011). The jet precession, in turn, can be attributed to (A) the presence of a companion, or (B) an asymmetry in the accretion of the envelope into the disk. The presence of a companion could lead to (a) the precession of the accretion disk due to the tidal interactions with a non-coplanar companion, or (b) the orbital motion of the jet source around a companion in orbit with the core. In case (B), the asymmetry could be due to a large difference in the orientation of magnetic and rotation axes of the core as demonstrated by simulations (Hirano & Machida 2019). ALMA studies at a resolution of 40 mas (∼56 au) did not reveal the presence of a binary companion within the MM1 core at these separations (Busquet et al. 2019), and therefore higher-resolution studies are required to draw valid conclusions regarding the presence of a binary companion.

The luminosities of knots in H₂ and [Fe ii] lines are calculated using dereddened fluxes and are in the range 10³¹–10³³ erg s⁻¹, tabulated in Table 3. For this, we have assumed A_v = 30 mag (Davis et al. 2001) for the knots close to the central region and in the northern arm (Knots 1, 2, 3, 4, and 5). Rodríguez-Kamenetzky et al. (2019) had calculated A_v = 2.8 mag and 3.4 mag for HH80 and HH81 (Knots 8 and 9), respectively. For the knots at intermediate distances (Knots 6 and 7), we have assumed A_v = 15 mag. We note that the ratio of [Fe ii] and H₂ luminosities are in the range 1.4–70.0. This is higher than the values obtained for low-mass class 0 and class I YSOs for which (Davis et al. 2003; Caratti o Garatti et al. 2006). We did not find this ratio for jets from massive YSOs in literature. We attempt to understand the observed ratio in terms of the dynamical age and evolutionary stage of the YSO driving the jet. The dynamical timescale of the HH80-81 radio jet is ∼10⁴ yr (Masqué et al. 2012), while the associated outflows provide an estimate of ∼10⁶ yr (Benedettini et al. 2004). As massive objects have shorter evolutionary timescales compared to their lower-mass counterparts, this would suggest that the driving source is at a later evolutionary phase when compared to the low-mass class 0/I sources of similar dynamical ages (Barsony & Kenyon 1992; Andre et al. 2000). This indicates that the difference in is plausibly the effect of evolutionary phase of the driving source, which manifests its influence on the jet physical properties. This is also corroborated by the large extent of the jet, ∼18 pc. Another noteworthy point is that the ratio implies a weakening of molecular line emission with the age of the driving source. This is expected because the clearing of the ambient medium by the previous mass-loss activities would result in lower densities and thus a weak ambient magnetic field (Hollenbach 1997), thereby favoring the dominance of dissociative J-type shocks over nondissociative C-type shocks.

Table 3. Physical Parameters of Jet Knots and Mass-loss and Momentum Rates for Knots Showing [Fe ii] Emission

Source Name		L[Fe II]	v⊥	l⊥
	(1031 erg s−1)	(1032 erg s−1)	(km s−1)	(″)	(10−5 M⊙ yr−1)	(10−2 M⊙ yr−1 km s−1)
Knot 1	8.9	⋯ a	⋯ a	⋯ a	⋯ a	⋯ a
Knot 2	6.8	⋯ a	⋯ a	⋯ a	⋯ a	⋯ a
Knot 3	13.2	13.7	1000	11.5	4.3	4.3
Knot 4	⋯ b	34.5	1000	24.0	5.2	5.2
Knot 5	2.0	13.6	1000	14.0	3.5	3.5
Knot 6	1.0	1.4	500	6.5	0.4	0.2
Knot 7	5.7	0.8	500	17.5	0.1	0.04
Knot 8	0.5	0.3	500	11.5	0.04	0.02
Knot 9	1.3	0.5	500	30.0	0.03	0.02

Notes.

^a [Fe ii] emission is not detected in these knots. ^b H₂ emission is not detected in this knot.

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We also compute the mass-loss rates () of the knots using the mass (M_jet), tangential velocity (v_⊥), and sky-projected length (l_⊥) by using the following formula

Here, μ = 1.24, m_H = 1.67 × 10⁻²⁷ kg, and A_i = 0.00465 s⁻¹ (Nussbaumer & Storey 1988) are mean atomic weight, proton mass, and radiative rate of the considered transition, respectively. We assume that the knot temperatures are ∼10⁴ K since radio emission is detected in the majority of our knots. For this temperature and electron densities of ∼10⁴ cm⁻³, which are typical of protostellar jets (Nisini et al. 2005; Takami et al. 2006; Giannini et al. 2013), the fractional population of the upper level of the transition f_i = 0.01 (Hamann et al. 1994) and (Hamann 1994). We assume a gas phase abundance of iron , which is the solar abundance (Asplund et al. 2021) assuming that there is no dust depletion.

To get estimates of the velocities of knots, we utilize proper-motion studies of HH80-81 from the literature. Radio observations have shown that the tangential velocities of knots close to the YSO can be as high as 1000 km s⁻¹ (Marti et al. 1993, 1995). An optical proper-motion study by Heathcote et al. (1998) had revealed a range of velocities for the subknots of HH81 and HH80, and here we adopt an average velocity of 500 km s⁻¹ for both these knots. We also assume v_⊥ = 500 km s⁻¹ for the knots at intermediate distances. The extent of [Fe ii] emission along the jet in each knot is adopted as the projected length (l_⊥).

The inferred mass-loss rates for the knots are in the range 3.0 × 10⁻⁷–5.2 × 10⁻⁵ M_⊙ yr⁻¹, tabulated in Table 3. For the knots close to the YSO, the values are consistent with those obtained from other tracers: 3 × 10⁻⁵ M_⊙ yr⁻¹ from molecular gas using CO (Qiu et al. 2019), and ∼10⁻⁵ M_⊙ yr⁻¹ using radio emission from ionized gas (Carrasco-González et al. 2012) toward the central region of the jet system. The mass-loss rates of HH80-81 jet are larger than those from jets of low-mass YSOs, 10⁻¹⁰–10⁻⁷ M_⊙ yr⁻¹ (Davis et al. 2003; Garcia Lopez et al. 2008; Dionatos et al. 2009; Reiter et al. 2016), but comparable to the [Fe ii] mass-loss rate from few massive YSOs, 10⁻⁷–10⁻⁴ M_⊙ yr⁻¹ (Caratti o Garatti et al. 2008; Fedriani et al. 2018, 2019). Assuming (Cabrit 2007; Antoniucci et al. 2008), we obtain of ∼10⁻⁶–10⁻⁴ M_⊙ yr⁻¹ for the knots. The accretion rates obtained from the knots close to the YSO are consistent with the observational and modeling estimates obtained toward the central region of this source (Carrasco-González et al. 2012; Añez-López et al. 2020) and is higher than the accretion rates observed toward low- and intermediate-mass protostars (Calvet et al. 2004; Evans et al. 2009; Ellerbroek et al. 2013; Antoniucci et al. 2014; Lee 2020).

In this section, we examine the region close to the driving source and its immediate surroundings. The driving source of the jet, MM1, is not visible in the NIR J, H, and K or the narrowband images, implying the highly embedded nature of the exciting YSO. A molecular outflow has been observed to be originating from MM1 with a position angle (PA) of 19° (Qiu et al. 2019) and is found to be associated with the radio jet. Molecular line observations using CO (2–1), (3–2), (6–5), and (7–6) lines by Qiu et al. (2019) have revealed the red- and blueshifted components of this outflow aligned along the direction of the radio jet. The outflow is marked as a white arrow pointing in the direction of the blueshifted emission in Figure 3. In the vicinity of MM1, we also observe two arc-shaped structures in the west and southwest that resemble bow shocks: arcs 8 and 9 in the figure. These arcs appear to surround MM1 toward the western side, and we believe that these could be bow shocks associated with the walls of wind-swept wide-angle cavities of strong winds (Davis et al. 2002) ejected by MM1.

Additional outflow components have been detected from the central region emanating toward northeast, northwest, and southeast using submillimeter SiO(2–1), HCO⁺(1–0), HCN(1–0), and CO(3–2) lines (Fernández-López et al. 2013). These outflows are shown as magenta and green arrows in Figure 3. The lengths of the arrows indicate the approximate extent of the detected emission from the outflows. These outflows are believed to originate from MM2, which is ∼10⁴ au from MM1. MM2 comprises two cores, toward the east MM2(E) and west MM2(W; Busquet et al. 2019). The string of H₂ knots in Reg 1 are aligned with the outflow in the southeast direction, and we believe that these H₂ knots are shock-excited by the associated jet. The extent of the knots matches the extent of the detected outflow lobe. We note that these five knots trace a curve, indicating precession of the outflow.

There is an H₂ knot in Reg 6 that appears to lie on the opposite side of this outflow. Although no outflow lobe has been detected in this direction, we believe that this knot could plausibly be due to the other lobe of the same outflow that gives rise to knots in Reg 1. Reg 4 has two H₂ knots that are aligned along the northeast SiO(2–1) outflow lobe. Again, we note a correspondence between the location of the knots and the extent of the outflow detected. Reg 2 has two H₂ knots that are located to the north of Reg 1 and south of Reg 4. It is difficult to comment on the likely outflow or exciting source responsible for this knot due to the cluster of sources present in the region. We also note arcs 11 and 12 facing inwards within the ellipse of Reg 7 appear to extend farther along the southeast of Reg 1. It is possible that they are bow shocks associated with the same outflow. Davis et al. (2001) had previously probed the central region (within 40″) using echelle spectroscopy and identified two bright H₂ knots and faint diffuse emission within the slit. The knots identified by them coincide with H₂ emission arcs that we discern toward the central region.

We also observe a number of H₂ arcs in Regs 3 and 5. These are seen to be surrounding the 8 μm nebulous regions that are located about ∼20″ from MM1. While it is possible that a few of these may be associated with outflows discussed above, we explore the possibility of YSOs in the vicinity as the source of excitation of these bow shocks. Of the several YSOs identified in this region by Qiu et al. (2008), we find that two YSOs, namely YSO1 and YSO2, are associated with MIR nebulosity as well as radio emission; see Figure 2. Figure 2 shows that Reg 5 envelopes the H₂ arcs 1–7 and 8 μm nebulosity around YSO1. That said, Reg 3 encloses H₂ arcs 10–13 and 8 μm nebulosity around YSO2. YSO1 and YSO2 have been identified as intermediate to massive YSOs by Qiu et al. (2008) based on NIR and MIR observations. We have identified these objects as Stage 0/I objects by constructing their spectral energy distributions and fitting them with the radiative transfer models of Robitaille et al. (2007). We thus believe that YSO1 and YSO2 could be driving strong winds that result in the formation of H₂ arcs in Reg 3 and 5, which surround these YSOs.