Photoionized Herbig–Haro Objects in the Orion Nebula through Deep High Spectral Resolution Spectroscopy. II. HH 204

Collimated matter jets and Herbig–Haro (HH) objects are phenomena associated with star formation (see Mundt & Fried 1983; Hartigan 1989; Reipurth & Bally 2001; Nisini et al.2005, and references therein). These objects are considered to have originated through a centrifugal-magnetic launch mechanism from young stellar objects (YSOs; see Schwartz 1983; Strom et al. 1983; Nisini et al. 2018, and references therein). These jets have a doubly important role: on the one hand, from their origin, they regulate the stellar accretion by removing the angular momentum, modifying the conditions of the matter of the disk (see Hartigan et al. 1994; Giannini et al. 2013, 2015, and references therein); on the other hand, as it passes through the surrounding medium, they modify the physical conditions of the environment.

Within the strong radiation field of the Orion Nebula, the HH objects immersed in it are photoionized, so the emission of the gas in photoionization equilibrium of the HH objects dominates the global emission over the thin cooling layer that is formed after the shock passage (Henney 2002). This makes it possible to study the chemical composition of the gas of these HH objects—which, in principle, must be the same as in the Orion Nebula—with standard methods for studying photoionized regions.

The HH object HH 204 is located in the central region of the Orion Nebula, just southeast of the Orion Bar, apparently close to the θ² Ori A star. It was observed by Münch & Wilson (1962) and classified as an HH by Cantó et al. (1980). The origin of the jet is usually associated with the Orion south molecular cloud (Orion-S; O'Dell et al. 2017a), an active star formation area of the Orion Nebula. However, the source of the driving jets that feed HH 204 is not entirely clear, as we discuss in this paper. Object HH 204 is photoionized by θ¹ Ori C from behind its direction of propagation through the cavity formed by the shock (O'Dell et al. 1997a, 2017a). Through long-slit spectra, Mesa-Delgado et al. (2008) studied the effects of HH 204 on the gas of the Orion Nebula, finding peaks in the density and temperature distributions when crossing its surrounding area, as well as increases in the emission flux of [Fe iii] lines produced by dust destruction. Using integral field spectroscopy, Núñez-Díaz et al. (2012) studied the influence of HH 204 in the Orion Nebula in an area of $16\times 16\,{\mathrm{arcsec}}^{2}$, finding the presence of a trapped ionization front, as well as arguments in favor of the location of the object within the main body of the Orion Nebula and not in the Veil. Work by Mesa-Delgado et al. (2008), Núñez-Díaz et al. (2012), and O'Dell et al. (2017a) shows the presence of a high-T_e([N ii]) zone, attributed to shock heating. However, this effect and the coincidental fall in the total abundance of O may be related to an underestimation of the electron density, n_e, an alternative explanation that will be discussed in Section 7.

This is the second paper in a series dedicated to studying photoionized HH objects in the Orion Nebula using high-resolution spectroscopy obtained with the Ultraviolet and Visual Echelle Spectrograph (UVES; D'Odorico et al. 2000) of the Very Large Telescope (VLT) and Hubble Space Telescope (HST) imaging. In this work, we analyze the physical conditions, chemical composition, and dynamical properties of HH 204, separating the emission of the Orion Nebula from the HH object and other ionized gas components present in the line of sight. Previous to the present paper, there were few works dedicated to high-resolution spectroscopy of photoionized HH objects of the Orion Nebula, such as HH 202 S (Mesa-Delgado et al. 2009), HH 529 II, and HH 529 III (Blagrave et al. 2006; Méndez-Delgado et al. 2021).

This paper has the following content. In Section 2, we describe the observational data and their treatment. In Section 3, we describe the measurement of spectral lines and the reddening correction. In Section 4, we derive the physical conditions and ionic abundances of each of the observed velocity components, while in Section 5, we focus exclusively on HH 204, deriving its physical conditions, ionic abundances, and some properties pixel by pixel along the UVES slit, as well as studying the spatial distribution of the emission of HH 204 with HST imaging. In Section 6, we estimate the total abundances of the observed gas components. In Section 7, we study the effects of mixing three gas components of very different densities along the line of sight, simulating a spectrum with lower spectral resolution. In Section 8, we investigate the origin of HH 204 and its relationship with HH 203. In Section 9, we discuss the main results of this work and their implications. Finally, in Section 10, we summarize the conclusions. In Appendix A, we show the reliability of the [Fe iii] atomic data that we use. In Appendix B, tables of data and figures are added as support material.

The observations were made during the nights of 2013 October 28 and 29 under photometric conditions using UVES in the UT2 of the VLT in Cerro Paranal, Chile. The slit position was centered at the coordinates R.A.(J2000) = 05^h35^m2272, decl.(J2000) = −05°25'2042 with a position angle (PA) of 137°. The slit width provides an effective spectral resolution λ/Δλ ≈ 6.5 km s⁻¹, covering the spectral range between 3100 and 10420 Å. Three exposures of 150 s of the standard star GD 71 (Moehler et al. 2014a, 2014b) were taken on the same night under similar observational conditions as the science images to achieve the flux calibration of the data. The observational settings are shown in Table 1, and the spatial coverage is presented in Figure 1. The instrumental configuration and data reduction procedure are described in Méndez-Delgado et al. (2021, hereafter Paper I). The 2D spectra (see Figure 2) show three evident components: (1) the nebular one (the emission of the Orion Nebula), which is rather homogeneously distributed along the spatial axis of the slit and occupies the reddest spectral position; (2) the "diffuse blue layer" (DBL), a slightly blueshifted homogeneous diffuse component (previously detected by Deharveng 1973) that may correspond to a different H ii region along the same line of sight (García-Díaz & Henney 2007); and (3) HH 204, the "ball-shaped" blueshifted component. We define two spatial cuts—shown in Figure 2—covering a spatial area of 738 for cut 1 and 197 for cut 2. In cut 2, we can separate the emission of the DBL and the nebular component. However, due to the strong contribution of HH 204, we cannot separate those components in cut 1. In this case, we study the emission of the combined spectrum of the nebular component and the DBL. We also take advantage of the quality of the data, performing a pixel-by-pixel analysis of various emission lines in order to detect small variations in physical conditions and/or the chemical composition of HH 204 along the slit.

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The study of the spatial distribution of the emission of HH 204 and the gas flows that may give rise to it are based in the HST WFPC2 imaging in the F502N $(\overline{{\boldsymbol{\lambda }}}=5012)$, F547M $(\overline{{\boldsymbol{\lambda }}}=5446)$, F656N $(\overline{{\boldsymbol{\lambda }}}=6564)$, and F658N $(\overline{{\boldsymbol{\lambda }}}=6591)$ filters from program GO 5469 (O'Dell & Wong 1996). The spatial pixel size of these data is 0045. Flux calibration and correction for contamination by continuum and nontarget lines was performed using the coefficients given in O'Dell (2009).

We use the SPLOT task from IRAF ⁶ (Tody 1993) to measure the line intensities and estimate their uncertainties, as described in detail in Paper I. In the case of the spectra of cuts 1 and 2, we measure a complete set of ∼500 and ∼300 emission lines, respectively, while in the case of the pixel-by-pixel measurements for HH 204, we limit the analysis to some representative lines: H9, Hβ, Hα; He i λ λ4471, 5876, 6678; [N ii] λ λ5755, 6584; O i λ7772; [O i] λ6300; [O ii] λ3726; [O iii] λ λ4363, 4959; [Ne iii] λ3869; [S ii] λ λ6716, 6731; [S iii] λ λ6312, 9531; [Cl ii] λ9124; [Cl iii] λ5538; [Ar iii] λ7136; [Ca ii] λ7324; [Cr ii] λ8000; [Fe ii] λ9052; [Fe iii] λ λ4658, 4702, 4881; [Ni ii] λ7378; and [Ni iii] λ7890. The reddening correction was done using the extinction curve from Blagrave et al. (2007) and the emissivity coefficients of Storey & Hummer (1995) for the Hε, Hδ, Hγ, Hβ, and Hα Balmer lines and the P12, P11, P10, and P9 Paschen lines. The values of the extinction coefficient, c(Hβ), are presented in Table 2. In the case of pixel-by-pixel measurements, a value of c(Hβ) = 0.42 ± 0.02 was used. An example of the spectra that can be found in the online material is shown in Table 9, where some lines of the spectra of cut 1 are shown.

4.1. Physical Conditions

We use version 1.1.13 of PyNeb (Luridiana et al. 2015) to obtain the physical conditions of the gas from the intensity ratios of collisionally excited lines (CELs) and recombination lines (RLs). PyNeb is a Python-based tool to compute line emissivities and derive physical conditions and chemical abundances of ionized gas. We have used the atomic data set presented in Tables 10 and 11 for the calculations made with PyNeb. We first estimate the n_e values given by each diagnostic of CELs by calculating each convergence of T_e − n_e with the available diagnostics of electron temperature, T_e, using the PyNeb task getCrossTemDen, as described in detail in Paper I. The density and temperature diagnostics used are shown in Table 3. Then, in the nebular and DBL components, we adopt the weighted mean ⁷ of the available values of n_e obtained with the following diagnostics: [O ii] λ3726/λ3729, [S ii] λ6731/λ6716, and [Cl iii] λ5538/λ5518. For consistency, in the case of HH 204, we rely on the n_e derived from the [Fe iii] lines, since values of 10⁴−10⁶ cm⁻³ are above the critical densities of the CELs involved in the more common diagnostics. The simultaneous estimation of n_e([Fe iii]) and T_e([Fe iii]) in HH 204 is achieved by a maximum-likelihood procedure, as described in Paper I. In this procedure, different combinations of T_e and n_e are tested to obtain the abundance of Fe²⁺/H⁺ with several [Fe iii] lines, giving as a result the combination of T_e−n_e that minimizes the dispersion between the abundances obtained with all of the lines. In HH 204, we have confident detections of [Fe iii] λ λ3240, 3335 lines from the ⁵D − ³D transitions, whose ratios with lines from the multiplets ⁵D − ³F and ⁵D − ³P are highly dependent on T_e, as shown in Figure 3. We include the following lines in the maximum-likelihood calculation: [Fe iii] λ λ3240, 3335, 4658, 4702, 4734, 4881, 5011, 5271. This collection of lines allows us to obtain well-constrained values of T_e([Fe iii]) and n_e([Fe iii]). The intensity ratios of these selected lines are consistent with the predicted ones when using transitions coming from the same atomic level (which are independent of the physical conditions of the gas), as we show in Table 12. Another density indicator that can be used with our data is n_e(O ii), but only for the nebular component, which is the only one where we detect RLs of multiplet 1 of O ii.

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Table 3. Physical Conditions Determined from Several Diagnostics

	Cut 1		Cut 2
Diagnostic	HH 204	Nebula + DBL	DBL	Nebula
ne ( cm−3)
[O ii] λ3726/λ3729	${\mathrm{15,420}}_{-3850}^{+7740}$	${1130}_{-110}^{+150}$	${400}_{-120}^{+140}$	${1480}_{-180}^{+190}$
[S ii] λ6731/λ6716	${\mathrm{11,350}}_{-3890}^{+9920}$	${1350}_{-260}^{+290}$	${300}_{-120}^{+140}$	${1230}_{-230}^{+250}$
[Cl iii] λ5538/λ5518	${\mathrm{13,370}}_{-1830}^{+1990}$	${1630}_{-320}^{+370}$	⋯	${1930}_{-650}^{+720}$
[Fe ii] λ9268/λ9052	${\mathrm{13,100}}_{-2990}^{+2860}$	⋯	⋯
[Fe iii] λ4658/λ4702	${\mathrm{13,040}}_{-3130}^{+3830}$	${3380}_{-1340}^{+1810}$	⋯	${3200}_{-1540}^{+2540}$
ne(O ii)	⋯	1350 ± 150	⋯	1050 ± 200
[Fe iii]*	13,540 ± 1210	⋯	⋯	⋯
Adopted	13,540 ± 1210	1230 ± 160	350 ± 50	1440 ± 170
Te (K)
${T}_{{\rm{e}}}\left({\rm{He}}\,{\rm{I}}\right)$	${8790}_{-430}^{+480}$	9760:	5650:	7980:
[O i] λ5577/λ λ 6300+64	${8290}_{-320}^{+430}$	⋯	⋯	⋯
[N ii] λ5755/λ6584	${8760}_{-180}^{+170}$	${8530}_{-190}^{+150}$	${8120}_{-360}^{+390}$	${8440}_{-210}^{+170}$
[O ii] λ λ3726+29/λ λ7319+20+30+31	⋯	⋯	${\mathrm{10,390}}_{-640}^{+730}$	${9120}_{-470}^{+430}$
[S ii] λ λ4069+76/λ λ6716+31	${8260}_{-500}^{+640}$	${\mathrm{11,470}}_{-630}^{+950}$	${\mathrm{10,440}}_{-1030}^{+1360}$	${9890}_{-610}^{+650}$
[O iii]λ4363/λ λ4959+5007	${\mathrm{12,430}}_{-220}^{+180}$	${8010}_{-80}^{+90}$	⋯	${8120}_{-100}^{+90}$
[S iii] λ6312/λ λ9069+9531	${9310}_{-330}^{+220}$	${8180}_{-230}^{+190}$	${7710}_{-400}^{+510}$	${8010}_{-210}^{+250}$
[Fe iii]*	8210 ± 220	⋯	⋯	⋯
Te (low) adopted	8760 ± 180	8530 ± 190	8120 ± 390	8440 ± 210
Te (high) adopted	12,430 ± 220	8030 ± 60	7710 ± 510	8110 ± 90

Note. An asterisk indicates that a maximum-likelihood method was used. The bold values are those adopted for the determination of chemical abundances.

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Once the representative n_e is adopted for each component, we estimate T_e through several diagnostics based on CELs, as shown in Table 3. In the case of T_e([S iii]), telluric absorptions affect the line λ9069 in the nebular and DBL components. Thus, we adopt I([S iii] 9531)/I([S iii] 9069) = 2.47 (Podobedova et al. 2009) in these cases. In HH 204, we were able to separate the auroral [O i] λ5577 line from sky emission contamination, which permitted us to estimate T_e([O i]). In the DBL, the estimations of T_e([O ii]) and T_e([S ii]) are affected by some extended residual emission of HH 204 in the auroral lines that crosses the cut border, affecting the first pixels of cut 2. The T_e(He i) was estimated using the average values obtained from He i I(λ7281)/I(λ6678), I(λ7281)/I(λ4922), and I(λ7281)/I(λ4388) line intensity ratios. Finally, we define T_e(low) as the weighted mean of T_e([N ii]), T_e([O ii]), and T_e([S ii]), while T_e(high) is the weighted mean of T_e([O iii]) and T_e([S iii]).

The resulting physical conditions for all components are shown in Table 3.

4.2. Ionic Abundances

4.3. Ionic Abundances of Fe and Ni Ions

In HH 204, the emission lines of [Fe ii], [Ni ii], [Fe iii], and [Ni iii] are considerably enhanced in comparison with what is observed in the nebular component. Due to the low ionization degree of HH 204, we expect that Fe⁺ and Ni⁺ have an important contribution to the total Fe and Ni abundances. Therefore, it seems pertinent to discuss in some detail the degree of confidence of the abundance determinations based on these two ions.

Optical lines coming from the upper levels of the Fe⁺ atom can be affected by continuum pumping (Lucy 1995; Rodríguez 1999; Verner et al. 2000). However, lower levels that produce the emission lines of multiplet a⁴F − a⁴P are mostly populated by collisions (Baldwin et al. 1996). One of the strongest lines of this multiplet, [Fe ii] λ8617 (a⁴F_9/2 − a⁴P_5/2), could not be detected due to the instrumental gap of UVES in the red arm. However, weaker lines arising from the same upper level as λ λ9052, 9399 (a⁴F_7/2 − a⁴P_5/2, a⁴F_5/2 − a⁴P_5/2), detected in HH 204, must be useful for the same purpose. Although the transition probabilities of the weakest detected lines coming from the a⁴P_1/2, a⁴P_3/2, and a⁴P_5/2 levels still need to be tested (since these lines may be affected by undetected telluric absorptions), there is a good agreement between the measured and predicted line ratios of [Fe ii] λ9052/λ9399, λ8892/λ9227, and λ9268/λ9034, as it is shown in Table 13. In order to make a simple test of the chosen atomic data, we take advantage of the theoretical density dependence between the population of the a⁴P_1/2 and a⁴P_5/2. By using the estimated T_e(low) for HH 204 and the [Fe ii] λ9268/λ9052 intensity ratio, we obtain ${n}_{{\rm{e}}}([{\rm{F}}{\rm{e}}\,{\rm{I}}{\rm{I}}])\,={13,100}_{-2990}^{+2860}\,{{\rm{c}}{\rm{m}}}^{-3}$, which is consistent with the rest of the density diagnostics shown in Table 3. In cut 2, we derive the Fe⁺ abundance of the nebular component by using the uncontaminated [Fe ii] λ8892 line.

The a⁴F − a⁴P_5/2 transitions of [Fe ii] and the a²D − a²F_7/2 ones of [Ni ii] have practically the same excitation energy, giving an origin for lines that are close in wavelength (Bautista et al. 1996). However, there is an important difference in their sensitivity to fluorescence by continuum pumping due to the multiplicity of their ground states. Photoexcitations from the Fe^{+ 6}D ground state to the quartet levels have low probability, and lines produced by intercombination transitions from sextet to quartet levels should be very weak (Bautista & Pradhan 1998). However, Rodríguez (1999) pointed out that the lowest quartet level, a⁴F_9/2, may be metastable and promote excitations to higher quartet levels. The main pumping routes starting from this level were studied by Verner et al. (2000) at densities above 10⁴ cm⁻³, finding that this pumping populates the levels a⁴H, b⁴F, b⁴P, and a⁴G. Since transitions from any of these levels to a⁴P are rather weak, its population remains practically unaffected. Nevertheless, in the case of [Ni ii], the ground state and participating levels are doublets that make fluorescence effects by continuum pumping more likely (Bautista et al. 1996). However, an important factor that plays against the influence of fluorescence effects in [Ni ii] in the case of HH 204 is its relatively large distance from θ¹ Ori C (1504), the main ionization source of the nebula (O'Dell et al. 2015, 2017b). In a simple procedure, following the formalism developed by Bautista et al. (1996, their Equation (8)), for a three-level model (level 1: a²D_5/2, level 2: a²F_7/2, and $\mathrm{level}\ 3:\ {{\rm{z}}}^{2}{{\rm{D}}}_{5/2}^{0}$), the critical densities n_cf—for which if n_e > n_cf, collisional excitations dominate over fluorescence—in two zones of the Orion Nebula (a and b), both excited by θ¹ Ori C, should be related as follows:

$\begin{eqnarray}&&\displaystyle \frac{{n}_{\mathrm{cf},\ {\rm{a}}}}{{n}_{\mathrm{cf},\ {\rm{b}}}}=\left(\displaystyle \frac{{J}_{13,\ {\rm{a}}}}{{J}_{13,\ {\rm{b}}}}\right)\left(\displaystyle \frac{{q}_{12,\ {\rm{b}}}}{{q}_{12,\ {\rm{a}}}}\right),\end{eqnarray} \tag{ 1 }$

where q₁₂ is the Maxwellian-averaged collisional strength for transitions from level 1 to level 2, and J₁₃ is the intensity of the continuum at energies of the 1 → 3 transitions. If we choose zone a as the one observed by Osterbrock et al. (1992) and zone b as HH 204, we can assume q_12,b/q_12,a ≈ 1, because the T_e determined by Osterbrock et al. (1992) and us are very similar (9000 and 8760 K, respectively). On the other hand, by estimating the geometrical dilution of J₁₃ in both areas (the zone observed by Osterbrock et al. 1992 is located 6398 from θ¹ Ori C), we get n_cf,a/n_cf,b ≈ 5.53. By adopting the n_cf,a estimated by Bautista et al. (1996), we obtain n_cf,b ≈ 2.17 × 10³ cm⁻³, which is rather small compared with the density we obtain for HH 204 and therefore collisional excitation should dominate. Nevertheless, it must be considered that the apparently closer star θ² Ori A may also be a source of fluorescence for HH 204. However, by using the [Ni ii] λ7378 (a²D_5/2 − a²F_7/2) line to obtain the Ni⁺ abundance and comparing with the Fe⁺/H⁺ ratio, we obtain $\mathrm{log}({\mathrm{Ni}}^{+}/{\mathrm{Fe}}^{+})=-1.27\pm 0.06$, which is in complete agreement with the solar value of $\mathrm{log}{(\mathrm{Ni}/\mathrm{Fe})}_{\odot }=-1.25\pm 0.05$ (Lodders 2019), suggesting the absence of significant fluorescence effects (as discussed before, we expect larger fluorescence effects in Ni⁺). Therefore, we can assume that θ² Ori A is not a significant source of photon pumping of [Ni ii] lines in HH 204. We do not estimate the Ni⁺ abundances for the rest of the velocity components because it requires a detailed analysis of the fluorescence conditions in the ionized gas, which goes beyond the scope of this paper.

We derive the Fe²⁺ abundance using the [Fe iii] lines indicated in Section 4.1. The good agreement is noticeable between T_e([Fe iii]), T_e([O i]), and T_e([S ii]) in the case of HH 204, contrary to what was found in HH 529 II and HH 529 III, where T_e([Fe iii]) was more consistent with the temperature obtained for high-ionization ions (Méndez-Delgado et al. 2021). This is not surprising due to the different ionization degrees of HH 204 and HH 529 II+III (see Section 5).

In Paper I, we pointed out the inconsistency between the predicted and measured intensity ratios of [Ni iii] ³F − ³P₂ transitions (λ λ6534, 6000, 6946) in HH 529 II, HH 529 III, HH 202 S, and several zones of the Orion Nebula (see Table D11 of Paper I). We obtain a similar result for HH 204, λ6534/λ6000 = 1.38 ± 0.18, which is rather far from the predicted value of 2.19 (Bautista 2001). This indicates that the transition probabilities of the aforementioned lines may have errors (for a more detailed discussion, see Appendix C in Méndez-Delgado et al. 2021). We have a different situation for the intensity ratios of lines arising from the ¹D₂ level. After subtracting the small contribution of [Cl iii] λ8499.60 to the measured intensity of [Ni iii] λ8499.62, we obtain [Ni iii] λ7890/λ8500 = 2.65 ± 0.19, in agreement with the predicted value of 2.47 (Bautista 2001). This indicates that, with the available atomic data, the most confident determinations of the Ni²⁺ abundance can be obtained with these last lines. Thus, we will adopt the Ni²⁺ abundances determined from the [Ni iii] λ7890 line. Unfortunately, this line is affected by a telluric emission feature in the nebular component; therefore, we have to rely on the [Ni iii] λ6534 line to determine the Ni²⁺ abundance for this component.

4.4. Ionic Abundances of Ca⁺ and Cr⁺

4.5. Ionic Abundances Based on RLs

For the nebular component, the He⁺ abundance is derived using T_e(high) and the lines considered in Table D14 of Paper I, which are the least affected ones by the metastability of the 2³S level. However, we have used T_e(low) for HH 204. In this component, our determination of T_e(He I) is more consistent with T_e(low). This is because in HH 204, [O iii] emission arises from a small localized area of higher-ionized gas, and T_e([O iii]) may not be representative of the He⁺ volume, as we describe in Section 5.1.

The C ii λ4267 is partially blended in the two velocity components of cut 1; therefore, we base our calculations on C ii λ9903. We use C ii λ4267 in cut 2. The C²⁺ abundance estimations based on both lines are in complete agreement in cut 2. Due to the similar ionization potentials of C⁺ and He⁰ and the considerations outlined in the previous paragraph, T_e(low) is also used for determining the C²⁺ abundance in HH 204.

Contrary to the situation presented in Paper I, in HH 204, O i RLs from multiplet 1 are severely affected by telluric emission features, with the exception of O i λ7772. We derive the O⁺ abundance of the HH object using the intensity of this line and the predicted line strengths from Wiese et al. (1996) following Equation (2) of Esteban et al. (1998).

Estimations of the O²⁺ abundance from RLs are based on the available O ii lines of multiplet 1. These are not detected in the case of HH 204 (see Figure 4). We use an estimate of the upper limit to the intensity of the λ4649 line for this component.

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As mentioned in Section 3, we measure several lines pixel by pixel along the slit. The spatial resolution in the blue and red arms of UVES is slightly different (0246 and 0182 pixel⁻¹, respectively). Cut 1 includes 30 pixels in the blue arm and 42 in the red one. In the pixel-by-pixel measurements, renormalization between lines in common in each arm is not enough to dilute possible differences in the integrated flux. However, Hβ is observed in the spectra of both arms; therefore, we split our pixel spectra into two parts, 27 blue pixel spectra and 37 red pixel spectra, both groups normalized with respect to F(Hβ). The missing first pixels (from east to west) of cut 1 of both arms were not included, since the emission of HH 204 was too faint. We proceeded as follows: based on the [Fe iii] λ4658/λ4702 line ratios, we derived n_e along HH 204 in the blue arm. Once the density distribution was estimated, the calculation of T_e([O iii]) was done, also in the blue arm, through the [O iii] λ4363/λ4959 line ratio. The spatial distribution of n_e was linearly interpolated in the red arm to estimate T_e([S iii]) and T_e([N ii]). Once the physical conditions are determined, we estimate the ionic abundances using the same procedure followed in Section 4.2. The zero-point of the spatial distribution is located at coordinates R.A.(J2000) = 05^h35^m2281, decl.(J2000) = −05°25'2186, just at the apparent eastern—external—edge of the bow shock. To estimate the distance from the bow shock along the jet, we adopt a heliocentric distance of 410 ± 10 pc (Binder & Povich 2018) to the Orion Nebula, based on Gaia DR2 parallaxes (Gaia Collaboration et al. 2018). The integrated emission is dominated by the blueshifted jet bullet component centered around ∼−20 km s⁻¹ (in a heliocentric velocity scale) within a 1σ range of ±10 km s⁻¹, being well separated from the DBL and the nebular emission.

5.1. Small-scale Physical Conditions

The resulting pixel-by-pixel distribution of physical conditions is shown in Figure 5. At the shock front, we can see that n_e([Fe iii]) reaches values up to a factor of about 2 higher than at a distance of ∼13 mpc from the bow shock. The distribution of T_e([N ii]) is practically constant, while T_e([S iii]) decreases slightly at the edge of the bow shock.

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Conversely, T_e([O iii]) strongly increases at distances closer to the bow shock as it is shown in Figure 6, where the intensity of the auroral line λ4363 is increased more than that of the nebular line λ4959. In the presence of a shock, a photoionized gas can be heated at a temperature higher than that fixed by photoionization equilibrium (Zel'dovich & Raizer 1967; see Section 11 of Paper I). After the shock passage, the gas cools down by radiative emission until reaching an equilibrium temperature, forming a cooling zone whose extension will be inversely proportional to the electron density (Hartigan et al. 1987). If we assume that the high-T_e([O iii]) area corresponds to the cooling zone formed after the shock, the fact that T_e([S iii]) and T_e([N ii]) are not affected in the same way suggests that the high ionization degree emission should come (at least partially) from a different gas volume than the one that originates the low-ionization emission. Therefore, we suggest that we are seeing the superposition of two different emission components: one from the bow shock and one from the Mach disk (the shock internal to the jet). This model will be discussed in Section 9.3.1.

5.2. Small-scale Patterns in the Ionic Abundances

Figure 7 shows the spatial distribution of the ionic abundances of O. As described in Section 5.1, the increase of T_e([O iii]) may be related to shock heating. Therefore, we highlight in red the O²⁺ abundances in this area in the bottom panel of Figure 7. In the top panel, we show the O⁺ abundances along the full distance range and the O ones in the area where T_e([O iii]) remains constant. This panel shows that practically all O is in O⁺ form. It should be noted that an increase of a factor of ∼2 in the O²⁺ abundance would represent less than 1% of the total O, well below the associated uncertainties; therefore, this increase would be undetected in analyses lacking our spatial and spectral resolutions.

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In Tables 4 and 5, we can see that the O⁺ abundances determined from CELs and RLs for HH 204 are the same within the errors, so we do not find an abundance discrepancy (AD) for this ion, contrary to the situation found in practically all photoionized nebulae. Figure 8 indicates the absence of systematical trends of the AD in the observed areas of HH 204. Although some fluctuations seem to be present, they are very small in any case.

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Table 4. Chemical Abundances Obtained with CELs of the Integrated Spectra of Each Component

	Cut 1		Cut 2
Ion	HH 204	Nebula + DBL	DBL	Nebula
O+	8.62 ± 0.05	8.14 ± 0.05	${8.26}_{-0.09}^{+0.13}$	${8.18}_{-0.05}^{+0.06}$
O2+	6.34 ± 0.02	7.96 ± 0.02	${7.33}_{-0.10}^{+0.15}$	8.04 ± 0.02
N+	7.72 ± 0.03	7.34 ± 0.03	${7.40}_{-0.06}^{+0.08}$	${7.29}_{-0.03}^{+0.04}$
Ne2+	5.05 ± 0.03	7.16 ± 0.02	⋯	${7.23}_{-0.02}^{+0.03}$
S+	6.60 ± 0.04	5.93 ± 0.03	${5.92}_{-0.06}^{+0.07}$	${5.86}_{-0.03}^{+0.04}$
S2+	6.80 ± 0.03	6.84 ± 0.03	${6.85}_{-0.08}^{+0.10}$	6.89 ± 0.04
Cl+	4.72 ± 0.03	4.17 ± 0.03	${4.08}_{-0.09}^{+0.10}$	4.05 ± 0.04
Cl2+	${4.77}_{-0.03}^{+0.04}$	4.93 ± 0.04	${4.99}_{-0.12}^{+0.16}$	${4.98}_{-0.05}^{+0.06}$
Ar2+	5.66 ± 0.03	6.10 ± 0.02	${5.99}_{-0.08}^{+0.10}$	6.12 ± 0.02
Ar3+	⋯	${3.64}_{-0.12}^{+0.13}$	⋯	⋯
Fe+	6.16 ± 0.04	⋯	⋯	4.72 ± 0.08
Fe2+	6.49 ± 0.02	5.72 ± 0.04	${5.56}_{-0.08}^{+0.10}$	5.77 ± 0.04
Fe3+	<5.11	5.73 ± 0.13	⋯	⋯
Ni+	4.89 ± 0.02	⋯	⋯	⋯
Ni2+	5.13 ± 0.03	4.37 ± 0.09	⋯	⋯
Ca+	3.50 ± 0.03	⋯	⋯	⋯
Cr+	4.28 ± 0.03	⋯	⋯	⋯

Note. Abundances in units of 12+log(Xⁿ⁺/H⁺).

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Table 5. Chemical Abundances Obtained with RLs of the Integrated Spectra of Each Component

	Cut 1		Cut 2
Ion	HH 204	Nebula + DBL	DBL	Nebula
He+	10.53 ± 0.02	10.85 ± 0.03	10.66 ± 0.06	10.92 ± 0.04
O+	8.57 ± 0.03	⋯	⋯	⋯
O2+	<7.54	8.25 ± 0.06	⋯	8.40 ± 0.03
C2+	7.76 ± 0.07	8.22 ± 0.04	⋯	8.37 ± 0.02

Note. Abundances in units of 12+log(Xⁿ⁺/H⁺).

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In Figure 9, we present the ionic abundances of Cl and S. The species of the same ionic stage of both elements show similar pixel-by-pixel distributions. The variations of the S²⁺/H⁺ and Cl²⁺/H⁺ ratios along HH 204 are comparatively much smaller than those of S⁺/H⁺ and Cl⁺/H⁺, which show a decrease of 0.8 dex along the diagram as the distance from the bow shock increases. At distances to the bow shock smaller than ∼4.9 mpc, the abundances of S⁺ and Cl⁺ seem to stabilize, and presumably, almost all S and Cl must be only once or twice ionized. This allows the estimation of their total abundances without an ionization correction factor (ICF).

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The pixel-by-pixel distributions of the ionic abundances of Fe and Ni are clearly correlated, as shown in Figure 10. Similar to those found for S and Cl (see Figure 9), close to the bow shock, the contribution of species of Fe and Ni with ionic charges higher than Fe²⁺ and Ni²⁺ to their total abundances should be negligible. The ratios of the ionic abundances between both elements remain constant, as shown in Figure 11, being $\mathrm{log}({\mathrm{Fe}}^{+}/{\mathrm{Ni}}^{+})=1.26\,\pm 0.03$, $\mathrm{log}({\mathrm{Fe}}^{2+}/{\mathrm{Ni}}^{2+})\,=1.37\pm 0.03$, and $\mathrm{log}(\mathrm{Fe}/\mathrm{Ni})\,=1.33\pm 0.03$. Although the value of $\mathrm{log}({\mathrm{Fe}}^{2+}/{\mathrm{Ni}}^{2+})$ is slightly above the recommended solar value ($\mathrm{log}{(\mathrm{Fe}/\mathrm{Ni})}_{\odot }=1.25\,\pm 0.05$; Lodders 2019), this may be the consequence of a slight systematic underestimation of Ni²⁺ abundance because, as we discussed in Section 4.3, the atomic data for this ion seem to show some inaccuracies.

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In Figure 12, we show the similar pixel-by-pixel distributions of the He⁺ and Ar²⁺ abundances. Both quantities decrease as we approach the bow shock due to the decrease of the ionization parameter as n_e increases. A slight increase is observed at distances less than ∼2.4 mpc, probably due to the same process discussed in Section 9.3.1 for the case of [O iii] lines. However, the impact of the shock contribution seems to be negligible for these ions. For example, the fact that T_e(He I) is consistent with T_e([N ii]) (see Section 4.1) reflects that the populations of the singlet levels, which are the ones used for determining T_e(He I), are largely unaffected.

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In Figure 13, we show that the abundance of N⁺ increases as we move toward the bow shock from $12+\mathrm{log}({{\rm{N}}}^{+}/{{\rm{H}}}^{+})\,=7.53\pm 0.03$ to an apparently constant value of 7.75 ± 0.02. That plateau indicates that all nitrogen should be only once ionized. Figures 14 and 15 show the pixel-by-pixel distributions of the Cr⁺ and Ca⁺ abundances, respectively, which are somewhat different from the ones of Fe⁺ or Ni⁺ (Figure 10). This indicates that the distributions of Fe⁺/Cr⁺ and Fe⁺/Ca⁺ ratios are not constant, contrary to what is obtained for Fe⁺/Ni⁺ (Figure 16). In the case of the Fe⁺/Cr⁺ ratio, the observed trend may be related to the slight differences between their ionization energies or different depletion patterns. The curve defined by the Fe⁺/Ca⁺ abundance ratio may be due to the coexistence of this ion and H⁰ in the trapped ionization front of HH 204 (see Section 5.3).

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5.3. Deuterium Lines in HH 204

Deuterium emission lines were first identified in the Orion Nebula by Hébrard et al. (2000a). Unlike the expected isotopic shift of −81.6 km s⁻¹ with respect to the hydrogen lines, they observed a shifted emission around ∼−71 km s⁻¹ from Hα and Hβ. The difference of ∼10 km s⁻¹ is essentially due to the fact that their emission is produced in different areas of the nebula, where the bulk of the gas is moving at different radial velocities. Since the hydrogen lines are produced by recombination in the ionized area that expands toward the observer, the deuterium emission is mainly due to fluorescence excitation by nonionizing far-UV continuum in areas slightly beyond the ionization front, such as the photon-dominated region (PDR), or in the H i–H ii interface (O'Dell et al. 2001). After the identification of deuterium emission lines in the Orion nebula, they were also identified in other H ii regions, such as M8, M16, DEM S 103, M20, and Sh 2-311 (Hébrard et al. 2000b; García-Rojas et al. 2005, 2006, 2007). As in the Orion Nebula, the deuterium emission in these H ii regions has a narrow line width, consistent with their origin in colder areas.

In this work, we detect the emission of Dζ, Dε, Dδ, Dγ, Dβ, and Dα, as shown in Figure 17. In Table 6, we present the characteristics of these emissions, including the radial velocity of the D i and H i lines with respect to the laboratory wavelength of the H i ones. The observed isotopic shift of −81.4 km s⁻¹ between the deuterium and hydrogen lines indicates that both kinds of lines arise from HH 204. The observed D i/H i intensity ratios are in good agreement with the predictions of the standard model developed by O'Dell et al. (2001) for the Orion Nebula, confirming the fluorescent nature of the D i emission. Considering that the emission of deuterium occurs in areas slightly beyond the ionization front, the detection of these lines implies that the ionization front must be trapped in HH 204, moving along with it, consistent with the results of Núñez-Díaz et al. (2012), as well as other evidence that will be discussed in Section 9.3.

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Table 6. Characteristics of Deuterium and Hydrogen Lines in HH 204

	D i		H i
λ0	vr a (km s−1)	FWHM (km s−1)	vr (km s−1)	FWHM (km s−1)	I(D i)/I(H i)×1000
3889.05 b	−103.34	13.80 ± 1.39	⋯	⋯	⋯
3970.07	−103.10	14.20 ± 3.40	−21.54	24.62 ± 0.02	2.99 ± 0.45
4101.73	−102.83	13.53 ± 0.86	−20.97	24.49 ± 0.01	2.24 ± 0.15
4340.46	−103.02	16.17 ± 1.40	−20.83	24.66 ± 0.01	2.10 ± 0.16
4861.32	−102.26	14.31 ± 1.90	−21.47	24.67 ± 0.01	1.06 ± 0.11
6562.80	−101.82	14.90 ± 0.79	−21.88	24.94 ± 0.01	0.58 ± 0.03

Notes.

^a With respect to the laboratory wavelength of the closest H i line (first column). ^b The H i λ3889.05 emission of HH 204 is blended with the nebular one of He i λ3888.65.

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5.4. Subarcsecond Imaging of HH 204

Figure 18(a) shows the ratio of surface brightnesses, R([O III]) = S([O III] λ5007)/S(Hα λ6563), calculated from HST WFPC2 observations in the F502N, F547M, F656N, and F658N filters from program GO 5469 (O'Dell & Wong 1996). It can be seen that the line ratio in the background nebula shows a pronounced gradient from R([O III]) ≈ 0.3 in the northeast to R([O III]) ≈ 0.5 in the southwest. ⁸ Inside the bow shock, the ratio is significantly smaller, for instance, falling from ≃0.4 to ≃0.2 along the length of the UVES slit.

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However, the most interesting feature of the R([O III]) image is the slight increase in the ratio that is seen in a thin layer along the leading edge of the bow shock. This is most clearly visible in the northern wing of HH 204, such as the area highlighted by a dotted box in the figure. The average profiles across the shock for this region are shown in Figure 18(b). The lower panel shows that the raw ratio (solid black line) increases only slightly above its value in the background nebula, which is because the brightness increase across the bow shock is only a small fraction of the background brightness, as can be appreciated in the upper panel. In order to isolate the emission of the shocked gas from that of the nebula, we calculate the background-subtracted line ratio,

$\begin{eqnarray}&&{R}^{{\prime} }([{\rm{O}}\,{\rm\small{III}}])=\displaystyle \frac{S([{\rm{O}}\,{\rm\small{III}}])-{S}_{\mathrm{BG}}([{\rm{O}}\,{\rm\small{III}}])}{S({\rm{H}}\alpha )-{S}_{\mathrm{BG}}({\rm{H}}\alpha )},\end{eqnarray} \tag{ 2 }$

under the assumption that S_BG for each line is constant along the profile. The result is shown as a gray histogram in the lower panel of the figure, which reveals a sharp peak of width ≈0.3 mpc that reaches a maximum value ${R}^{{\prime} }([{\rm{O}}\,{\rm\small{III}}])\approx 2{R}_{\mathrm{BG}}([{\rm{O}}\,{\rm\small{III}}])$ and is centered on a displacement of ≈−0.1 mpc. The origin of the displacement axis is set to the peak in the spatial gradient of the Hα surface brightness, corresponding to the outer edge of the dense shocked shell. The negative displacement of the ${R}^{{\prime} }([{\rm{O}}\,{\rm\small{III}}])$ peak means that this occurs outside the dense shell, closer to the shock front itself.

Figure 18(c) shows the same quantities calculated along a cut that coincides with our UVES slit at the head of HH 204. In this case, ${R}^{{\prime} }([{\rm{O}}\,{\rm\small{III}}])$ is always significantly less than R_BG([O III]), but it does still show a small local peak with a position and width that are similar to the more impressive ones in the northern wing. These peaks in ${R}^{{\prime} }([{\rm{O}}\,{\rm\small{III}}])$ occur over a much smaller scale than any of the spatial gradients that we find in our UVES slit spectra and are only detectable because of the high spatial resolution of the HST. ⁹ For example, the increase in T_e([O III]) that we detect in the blueshifted emission near the shock front (Figure 5) occurs over a scale of 5 mpc, indicated by the red arrow in the figure, which is more than 10 times larger than the width of the ${R}^{{\prime} }([{\rm{O}}\,{\rm\small{III}}])$ peak.

What is the origin of the narrow peak in the [O III]/Hα ratio that is seen just outside the shocked shell? When a shock propagates into low-ionization gas (predominantly O⁺), there are three zones where enhanced [O III] emission might be expected (Cox & Raymond 1985; Sutherland & Dopita 2017): first, the radiative precursor in the preshock gas; second, the nonequilibrium collisional ionization zone immediately after the shock; and third, the radiative relaxation zone, where the postshock gas cools back down to the photoionization equilibrium temperature of ∼10⁴ K. The first of these can be ruled out in the case of HH 204 because the preshock photoionization of O⁺ would require shock velocities greater than 150 km s⁻¹, and observed proper motion and radial velocities imply a shock velocity around 100 km s⁻¹. The second zone has a high temperature (>50,000 K for shock velocities >55 km s⁻¹) but is severely underionized, resulting in line emissivities that are far in excess of the equilibrium values in a very thin layer. The third zone, in which oxygen is recombining through the O⁺⁺ stage while cooling through the range 30,000–10,000 K, is predicted to be somewhat thicker and with a higher electron density, yielding a greater contribution to the total [O III] emission. Given the electron density that we derive of 13,540 cm⁻³ (Table 3), and assuming a shock velocity <70 km s⁻¹, the cooling length should be approximately 0.05 mpc, or 0025, which is a few times smaller than the HST resolution. However, this analysis applies only to the head of the bow shock. In the wings, the shock is not perpendicular to the upstream gas velocity but is oblique at an angle α. This yields a postshock equilibrium density that is smaller by a factor of ${\cos }^{2}\alpha$ and a cooling length that is larger by the same factor. Hence, the cooling length is expected to be resolved for α smaller than about 45°, which is consistent with our observations of the narrow peak in the [O III]/Hα ratio in the north wing. The reason that the same behavior is not seen in the opposite wing is probably that the ambient nebular emission is much more highly ionized there, which masks the effect.

In the case of the nebular and DBL components, total abundances of O, Cl, and S were estimated by simply adding the abundances of all of their observed ions. Although there may be some contribution of S³⁺ and Cl³⁺, the ICFs of Stasińska (1978) and Esteban et al. (2015), respectively, predict negligible amounts of those species. In the case of N, Ne, Ar, and C, we adopt the same ICFs used by Arellano-Córdova et al. (2020). For Fe, we use the two ICFs proposed by Rodríguez & Rubin (2005). Since the real value of Fe should be between the predictions of both ICFs (Rodríguez & Rubin 2005), in Table 7, we present those determinations as lower and upper limits of the Fe abundance.

Table 7. Total Abundances

	Cut 1		Cut 2
Element	HH 204	Nebula + DBL	DBL	Nebula
O	8.62 ± 0.05	8.36 ± 0.03	8.31 ± 0.12	8.42 ± 0.04
O a	8.57 ± 0.03	⋯	⋯	⋯
N	7.75 ± 0.02	${7.56}_{-0.03}^{+0.04}$	${7.45}_{-0.08}^{+0.09}$	7.53 ± 0.05
Ne	⋯	7.56 ± 0.04	⋯	7.61 ± 0.05
S	7.07 ± 0.03	6.90 ± 0.03	6.90 ± 0.09	6.94 ± 0.04
Cl	5.10 ± 0.04	5.00 ± 0.03	5.04 ± 0.14	5.03 ± 0.05
Ar	⋯	6.14 ± 0.02	6.09 ± 0.10	6.17 ± 0.02
Fe	6.67 ± 0.03	5.91–6.09	5.64–6.19	5.97–6.13
Ni	5.35 ± 0.03	⋯	⋯	⋯
C a	⋯	8.49 ± 0.05	⋯	8.64 ± 0.04

Notes. Abundances in units of 12+log(X/H).

^a Based on RLs.

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In the case of HH 204, based on the results of Section 5, we decided not to derive the total abundances of elements for which we only observe highly ionized ions, such as He, Ne, Ar, and C, due to the low ionization degree of the gas and the large contribution of the ICFs. In the cases of O, N, Cl, S, Fe, and Ni, we can determine their total abundances without ICFs. As seen in Section 5, the spatial distribution of the abundances of the once- and twice-ionized ions of Cl, S, Fe, and Ni reach constant values at positions close to the bow shock, where the degree of ionization becomes very low. In this zone, the contribution of three or more times ionized ions of these elements should be negligible. A similar situation occurs with N, where the contribution of N²⁺ is expected to be very small close to the bow shock. Therefore, in Table 7, we present the total abundances obtained by adding the mean abundances of the once- and twice-ionized ions of Cl, S, Fe, and Ni for distances less than 4.9 mpc from the bow shock. In the case of O and N, we only consider the abundance of once-ionized ions in the same range of distances. At these distances, the pixel-by-pixel values of the O⁺ abundance determined from RLs have large errors (see Figure 8) because of the faintness of the O i λ7772 line. In this case, we use the O⁺ abundance obtained from the integrated spectrum presented in Table 4 to determine the total O abundance based on RLs.

In this section, we simulate a spectrum with lower spatial and spectral resolution by adding the flux of all of the velocity components, which includes the emission of HH 204, the DBL, and the emission of the Orion Nebula along the whole UVES slit. Following the reddening correction procedure described in Section 3, we obtain c(Hβ) = 0.36 ± 0.02 for this integrated spectrum.

In Figure 19, we present the resulting plasma diagnostics of the low-resolution spectrum. This diagram can be compared with those of the individual components, shown in Figure 25. If one only has the information provided by this degraded spectrum and applies the classic procedure of averaging n_e([O II]), n_e([S II]), and n_e([Cl iii])—excluding n_e([Fe iii]), since the sometimes-discrepant values given by this diagnostic are generally interpreted as the effect of incorrect atomic data—we would obtain n_e = 3430 ± 580. Using this value of density, we would obtain ${T}_{{\rm{e}}}([{\rm{O}}\,{\rm\small{II}}])\,={12140}_{-930}^{+950}$, ${T}_{{\rm{e}}}([{\rm{S}}\,{\rm\small{II}}])\,={19220}_{-2530}^{+9020}$, T_e([N II]) = 9200 ± 200, ${T}_{{\rm{e}}}([{\rm{S}}\,{\rm\small{III}}])={8740}_{-200}^{+230}$, and ${T}_{{\rm{e}}}([{\rm{O}}\,{\rm\small{III}}])={8530}_{-120}^{+100}$ K. It must be noted that the resulting T_e([N ii]) is higher than the ones obtained for each individual component analyzed in Section 4.1. Moreover, T_e([O II]) and T_e([S II]), the most density-dependent diagnostics, show much higher values. However, their effect on abundance determinations could be somehow mitigated, as their associated uncertainties are very high, and the use of a weighted mean of the different temperature indicators would reduce their contribution. The T_e([N ii]) always has much lower uncertainties and is generally the preferred temperature diagnosis for low ionization degree ions.

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Following the usual procedure and assuming the physical conditions determined in the previous paragraph, we would determine the O⁺ abundance using T_e([N ii]) and the O²⁺ one with T_e([O iii]), obtaining O⁺ = 8.15 ± 0.04, O²⁺ = 7.63 ± 0.02, and a total abundance of O = 8.26 ± 0.03. This value of the O/H ratio is lower than the one determined for all of the individual components. The only exception could be the DBL in cut 2, which shows O = 8.31 ± 0.12 (see Table 7), whose uncertainty is large enough to encompass the value obtained for the low-resolution spectrum. However, this does not mean that the DBL dominates the observed abundance of O, since it is the weakest component. This is demonstrated in Figure 20, which shows the line profile of f([O ii] λ3727), one of the most intense lines in the spectrum of the DBL.

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The low O abundance obtained in the low-resolution spectrum is due to the use of the classical diagnostics to estimate n_e, which do not adequately account for the high density of HH 204. The critical densities of the levels involved in those diagnostics are below the density of HH 204 (see Table D5 of Paper I). Likewise, the sensitivity of I([O II] λ3726)/I([O II] λ3729) and I([S II] λ6731)/I([S II] λ6716) at n_e ∼ 10⁴ cm⁻³ is much lower than at n_e ∼ 10²–10³ cm⁻³, the normal range of densities in H ii regions. The degree of ionization of each component also plays an important role. Although I([Cl III] λ5538)/I([Cl III] λ5518) is more density-sensitive than I([O II] λ3726)/I([O II] λ3729) or I([S II] λ6731)/I([S II] λ6716) at densities of around n_e ∼ 10⁴ cm⁻³, HH 204—the component with the highest density—has a very low degree of ionization. Therefore, in the combined emission of HH 204 and the Orion Nebula, the last component has a greater weight in n_e([Cl iii]). On the other hand, I([Fe III] λ4658)/I([Fe III] λ4702) is practically insensitive at densities smaller than n_e ∼ 10³ cm⁻³, and the critical density of this diagnostic is above ∼10⁶ cm⁻³. In addition, most of the [Fe iii] emission comes from HH 204 due to its higher abundance of gaseous Fe with respect to the Orion Nebula and DBL. These properties make it an excellent indicator of the presence of high-density gas as in HH objects. In our case, the ${n}_{{\rm{e}}}([{\rm{F}}{\rm{e}}\,{\rm{I}}{\rm{I}}{\rm{I}}])={10,790}_{-2620}^{+3230}$ cm⁻³ we obtain for the low-resolution spectrum is rather close to the density of HH 204. This confirms the importance of the warning given by Morisset (2017), who, through photoionization models, predicted large errors in the determination of the physical conditions and chemical abundances in nebulae if one assumes a single component when, in fact, there are several, and some of them are composed of high-density gas. The exercise we present in this section is an observational confirmation.

If instead of using the classical diagnostics to determine n_e, we take the average of the densities obtained for each component (see Table 3) weighted by their observed F(Hβ), we get n_e = 6820 ± 810 cm⁻³. This value is roughly between the predictions of classical diagnostics and n_e([Fe iii]). Note that in Figure 19, close to this value of density, T_e([O II]) and T_e([S II]) converge to T_e([N ii]). Using that density, we obtain ${T}_{{\rm{e}}}([{\rm{O}}\,{\rm\small{II}}])={8650}_{-520}^{+410}$, ${T}_{{\rm{e}}}([{\rm{S}}\,{\rm\small{II}}])={9890}_{-990}^{+1100}$, ${T}_{{\rm{e}}}([{\rm{N}}\,{\rm\small{II}}])\,={8850}_{-180}^{+210}$, ${T}_{{\rm{e}}}([{\rm{S}}\,{\rm\small{III}}])\,={8800}_{-160}^{+250}$, and ${T}_{{\rm{e}}}([{\rm{O}}\,{\rm\small{III}}])={8490}_{-120}^{+90}$ K. Calculating the ionic abundances of oxygen with these physical conditions, we obtain ${{\rm{O}}}^{+}={8.36}_{-0.05}^{+0.06}$ and O²⁺ = 7.64 ± 0.02, which implies O = 8.44 ± 0.05. These values are more consistent with those obtained in the analysis of the individual components.

It is clear that the discrepancy between the different density diagnostics is not necessarily an artifact of the atomic data used. Instead, each diagnostic may be revealing the changing conditions of the gas along the line of sight of the spectrum differently. Relying uncritically only on those density diagnostics that are consistent with each other could lead to significant systematic errors.

At least two different high-velocity flows converge on the general HH 203/204 region from the direction of the inner Orion Nebula (see Figure 21), but it is not clear if either of them are directly responsible for driving the HH 204 bow shock. One flow is at a PA of ≈118° and transitions from a high-ionization state northwest of the Bright Bar (cyan contours in Figure 21(a)) to a lower-ionization state (yellow contours) to the southeast of the Bright Bar. The other is at PA ≈ 140° and is of low ionization for its entire detected length. Both of these flows give the appearance of driving HH 203, which implies that HH 203 may be a superposition of two unrelated bow shocks. Such a superposition is consistent with the detection of two different velocity components (−73 and −39 km s⁻¹) at the head of the bow shock and also with the complex structure apparent in high-resolution HST images (see Figure 21(b)). O'Dell et al. (2015) noted that in addition to the main bow shock (HH 203a), there appears to be a second faint bow shock (HH 203b), associated with the PA 118 flow. We also detect a third faint bow shock, which we denote HH 203c, situated in front (southwest) of HH 204a. Note that O'Dell et al. (2015) gave PAs of 124° and 127°, respectively, for HH 203 and HH 204, which probably represent an average of the PA 118 and PA 140 flows.

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The southern portion of HH 203a, which we label as "[S ii] knot" in the figure, is particularly strong in the [S ii] and [O i] filters and coincides with the peak of the −39 km s⁻¹ feature. The spatial alignment and the similarity in velocity and ionization make it likely that this knot is part of the PA 140 flow. It is conceivable that this flow may extend farther to the southwest and be driving the HH 204 bow shock, although there is no direct evidence for this. On the other hand, a third flow at PA ≈ 108° is seen to feed into HH 204 from the west. This jet, first noted by Doi et al. (2004), is very short and stubby and can be traced back only 10'' (20 mpc) from the bow shock. There is another faint filament of high-velocity [O III] emission that extends between the HH 203 and HH 204 regions at PA ≈ 108° (see Figure 21(a)). This appears to provide a connecting bridge between the PA 140 and PA 108 flows, although the difference in velocity and ionization with respect to the PA 108 flow argues against a physical association with HH 204.

We have searched archival observations in other wave bands for any evidence of jets along the back-projection of the PA 108 axis. The most convincing association is with a molecular hydrogen filament seen in the 2.12 μm line (see Figure 21(e)). At the position of this filament, HH 204 is at PA = 105°, which is well within the uncertainties, and the orientation of the filament is consistent with the same PA. Unfortunately, no kinematic observations are currently available for this filament, so its association with HH 204 can only be tentative. The stellar source that best aligns with the H₂ filament is COUP 484; see Figure 21(e). However, this is a rather low-luminosity star and therefore seems an unlikely candidate for driving such an impressive large-scale outflow. The star V2235 Ori is also marginally consistent within the uncertainty with the PA 108 axis and is roughly 100 times brighter than COUP 484 in the K and L infrared bands (Muench et al. 2002), but its position is completely inconsistent with being the source of the H₂ filament. There is also marginal evidence from MUSE observations (Weilbacher et al. 2015) for a blueshifted [Fe iii] filament that extends from the position of the H₂ filament toward HH 204, but the data are noisy.

A further important line of evidence for the flow direction is provided by proper-motion measurements. We have remeasured the proper motions using HST images over an interval of 19 yr (1996–2015) using the methodology described in Section 1 of Paper I. For the "nose" of the HH 204 bow shock, we find a plane-of-sky velocity of (71 ± 9) km s⁻¹ at PA = (136 ± 3)°. After correcting to a common distance of 417 pc, the previous measurements of Doi et al. (2002) are (83 ± 10) km s⁻¹ at PA = (137 ± 7)°, which are consistent with our measurements within the uncertainties. The proper-motion axis is shown by a large red arrow in Figure 21(b) for comparison with the candidate axes from the high radial velocity jets. It is marginally consistent with the PA 140 axis but not with the PA 108, PA 118, or PA 164 axes.

In summary, convincing evidence for which large-scale flow might be driving the HH 204 bow shock is frustratingly absent. Although the PA 108 flow is clearly associated with HH 204, its short length means that the exact orientation is very uncertain. The PA 140 flow has a much better defined direction, but its extension beyond the position of the [S ii] knot in order to feed into the HH 204 bow shock is purely speculative. However, the close agreement between this flow direction and the proper-motion axis is an additional argument in its favor. The only thing that can be said with any degree of certainty is that the high-ionization PA 118 flow is not driving HH 204, only HH 203.

In Figure 21(b), we show the back-projection of all three of these flows into the core of the nebula, assuming an uncertainty of ±10° for the PA 108 flow and ±5° for the other two. The PA 118 flow is consistent with an origin in the Orion-S star-forming region, as has been remarked many times previously (O'Dell et al. 1997b; Rosado et al. 2002; O'Dell & Doi 2003). However, neither of the other flows are consistent with an origin in that region, unless the flow has suffered a relatively large-angle deviation. The back-projection of PA 108 falls significantly to the south of the main Orion-S region in an area with no convincing candidates for the driving source (see above discussion of the possible H₂ jet). The back-projection of the PA 140 flow intersects the Trapezium stars in the very center of the nebula, which raises the possibility that the source may be a proplyd, which are highly concentrated in that region.

The high spectral resolution of our data (λ/Δλ ≈ 6.5 km s⁻¹) allows us to identify and properly separate three kinematical components of ionized gas: the DBL, the emission of the Orion Nebula, and HH 204. In the following, we will discuss in detail the results concerning each of these components.

9.1. The DBL

The component designated as the DBL was first reported by Deharveng (1973), although it has been little studied, since high spectral resolution is required to separate its emission from that of the Orion Nebula. García-Díaz & Henney (2007) analyzed the velocity structure of the Orion Nebula through the emission of [O i], [S ii], and [S iii] lines using echelle spectroscopy. They detected the emission of the [S ii] doublet from the DBL, estimating a density of ∼400 cm⁻³, which is in complete agreement with our estimates. These authors did not detect the emission of [O i] or [S iii] in this component, although the emission of other low-ionization ions such as [O ii] and [N ii] was detected in previous works (Jones 1992; Henney & O'Dell 1999). This limited spectroscopic evidence led to interpreting the DBL as composed of fully ionized gas whose ionizing radiation field was rather soft, probably coming from θ² Ori A. We have detected all of these lines, along with [O i] and [S iii] ones, in the spectrum of this component extracted from cut 2 (see upper and middle panels of Figure 22). These emissions were also reported by O'Dell (2018) in a later reanalysis of the atlas of lines of García-Díaz & Henney (2007). In addition, we detect a weak [O iii] emission, indicative of the presence of gas with a high degree of ionization, as shown in the lower panel of Figure 22.

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Through observations of H i 21 cm emission, van der Werf et al. (2013) determined the existence of several H i velocity components in the Orion Nebula. In the southeast, in the area where the DBL is located, these authors identified a blueshifted component named "D," interpreted as an expanding shell centered on θ² Ori B, which is consistent with a scenario where this star ionizes the DBL. The observed [O i] emission is consistent with the presence of an ionization front in this nebular feature. However, with the new information provided by the ionic abundances of the DBL—estimated for the first time in this work—the simple model where the gas is photoionized exclusively by θ² Ori B may not be correct. Although small, the contribution of O²⁺ to the total abundance is not negligible, being around 10%. On the other hand, assuming that the DBL should have a chemical composition similar to the Orion Nebula, this implies that the estimated N⁺ abundance is approximately 75% of the total nitrogen abundance; therefore ,N²⁺ should be present in this component. Since θ² Ori B is a B0.7V star (Simón-Díaz 2010), we do not expect such a star to emit a number of photons capable of maintaining a significant proportion of highly ionized ions. This is reinforced by the spectroscopic results of Galactic H ii regions ionized by B-type stars such as Sh 2-32, Sh 2-47, Sh 2-82, Sh 2-175, Sh 2-219, Sh 2-270, Sh 2-285, Sh 2-297, and IC 5146 (García-Rojas et al. 2014; Esteban & García-Rojas 2018; Arellano-Córdova et al. 2021). In all of these regions, nitrogen is only once ionized, and the contribution of O²⁺ to the total oxygen is lower than 2%, with the exception of the faint Sh 2-47, although the O²⁺ abundance determination in this object is very uncertain.

As we can see in the discussion above, the spectroscopic properties of the DBL suggest some ionization by radiation leakage from the Orion Nebula. Simón-Díaz et al. (2011) found abnormal emission of CELs of highly ionized species (mainly [O iii]) in the external zones of M43, an H ii region ionized by a B0.5V star located to the northeast of the Orion Nebula. As those authors demonstrated, the spectral properties of this abnormal emission are consistent with contamination by scattered light from the Huygens Region. In our case, we can discard the scattered nature of the emission of highly ionized ions in the DBL because (i) it has the same velocity as the lines of low-ionization ions, and (ii) we do not detect anomalies in the Balmer decrement of the spectrum of the DBL, which would be a signature of the presence of scattered emission (see Simón-Díaz et al. 2011). Further observations with a longer exposure time, similar spectral resolution, and covering different areas of the Orion Nebula would shed more light on the extension and physical, chemical, and geometrical properties of the DBL.

9.2. The Nebular Component

9.3. HH 204

9.3.1. Two-zone Model for Observed Temperature Structure

Our spectroscopic observations allow us to analyze the physical conditions and ionic abundances of HH 204 with unprecedented detail. As shown in Section 5.1, the gas density is higher near the bow shock. On the other hand, only T_e([O iii]) seems affected by the shock, while T_e([N ii]) and T_e([S iii]) maintain their photoionization equilibrium values. This result may be explained by the weakness of high-ionization emission from the densest postshock gas in the bow shock and jet, allowing a greater relative contribution of the immediate postshock cooling zone to the [O iii] lines.

At each position along the spectrograph slit, the line of sight will cross several zones with different physical conditions, as illustrated in Figure 23:

• A1: the compressed shell behind the bow shock, which is in photoionization equilibrium;
• A2: the main body of the jet bullet, also in photoionization equilibrium;
• B1: the immediate postshock cooling zone of the bow shock; and
• B2: the postshock cooling zone of the jet shock.

In HH 204, the relative velocity between the unshocked jet and the working surface is very low (≈15 km s⁻¹), so the jet shock is much weaker than the bow shock, implying that the emission from zone B2 can be neglected compared with B1. Zones A1 and A2 should have similar conditions and so can be merged into a single zone with density n_A and temperature T_A. Although zone B1 should have a range of temperatures, for simplicity, we assume a single characteristic temperature T_B. The density of zone B is found by assuming a pressure equilibrium with zone A: n_B = n_A T_A/T_B. We define f_B for a given ion as the fraction of the total ionic emission measure, ∫n_e n_ion dz, that comes from zone B, with the remainder, f_A = 1 − f_B, coming from zone A.

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The appropriate value of T_B is rather uncertain, since it depends on the nonequilibrium evolution of ionization and temperature in the postshock radiative relaxation layer. Most published shock models (Cox & Raymond 1985; Sutherland & Dopita 2017) are calculated on the assumption that the far upstream and downstream ionization states are determined by the radiation from the shock itself. Care must therefore be exercised when translating their results to cases such as HH 204, where external irradiation from O stars is a dominant factor. The curved bow shock in HH 204 should give a range of shock velocities up to a maximum of V ≈ 84 km s⁻¹ (assuming the preshock medium is stationary). In principle, this corresponds to postshock temperatures as high as 2 × 10⁵ K, but the gas at such temperatures will be too highly ionized to significantly emit optical lines. The cooling timescale is generally shorter than the recombination timescale, so the gas is overionized as it cools. It is only when the temperature falls below about 50,000 K that the abundance of O⁺⁺ becomes significant (e.g., Figure 11 of Allen et al. 2008), allowing the emission of the optical [O iii] lines. A similar situation is seen in middle-aged supernova remnants, such as the Cygnus Loop (Raymond et al.2020).

We look for solutions where both T_A and T_B are constant along the slit, so that any spatial variation in the temperature diagnostics is driven primarily by variation in f_B. Although the density diagnostics do show a gradient with position, both T([O III]) and T([S III]) are relatively insensitive to density, so for simplicity, we assume n_A is constant. We use the Python library PyNeb to calculate the per-zone emission coefficients, j(T_A, n_A) and j(T_B, n_B), for each emission line. For a given diagnostic line pair, 1 and 2, the ratio is calculated as

$\begin{eqnarray}&&{R}_{12}=\displaystyle \frac{(1-{f}_{{\rm{B}}})\,{j}_{1}({T}_{{\rm{A}}},{n}_{{\rm{A}}})+{f}_{{\rm{B}}}{j}_{1}({T}_{{\rm{B}}},{n}_{{\rm{B}}})}{(1-{f}_{{\rm{B}}})\,{j}_{2}({T}_{{\rm{A}}},{n}_{{\rm{A}}})+{f}_{{\rm{B}}}{j}_{2}({T}_{{\rm{B}}},{n}_{{\rm{B}}})}.\end{eqnarray} \tag{ 3 }$

This is then fed into PyNeb's getTemDen function to find the equivalent single-zone temperature that would give the same ratio (assuming a density of n_A). It is clear from Equation (3) that for f_B = 0, one must recover T_e = T_A, and that for f_B = 1, one must recover T_e = T_B. But for intermediate values of f_B, the derived temperature will differ between ions because of variations in the temperature sensitivity of the diagnostic ratios.

We first investigate the case of a common f_B for all ions, but we find that this is unable to reproduce the observations. This is demonstrated in Figure 24(a), which shows the relation between T_e([O III]) and T_e([S III]) for four different values of T_B between 15,000 and 50,000 K. We set T_A = 9000 K and n_A = 20,000 cm⁻³ in all cases, and f_B increases from left to right along each curve. The gray rectangle shows the observed range of temperatures along the spectrograph slit (Figure 5); T_e([O III]) shows a systematic decline from ≈17,000 K near the bow shock to ≈12,000 K further away, while T_e([S III]) is roughly constant at 9000–10,000 K, with no apparent correlation with T_e([O III]). The two-zone models with T_B ≥ 30,000 K all show T_e([O III]) > T_e([S III]) as f_B increases, but this is insufficient to explain the observations. For example, in order to achieve T_e([O III]) = 17,000 K, the models predict T_e([S III]) > 11,000 K, which is significantly higher than observed.

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In Figure 24(b), we relax the assumption of a common f_B for all ions, separately showing the predicted values of T_e([N II]), T_e([S III]), and T_e([O III]) as a function of f_B, assuming T_B = 30,000 K. The ranges of observed values are shown by colored bands, blue for [O III] and orange for [N II] and [S III]. From the figure, it is apparent that a decline from f_B([O III]) ≈ 0.1 at x = 0 to f_B([O III]) ≈ 0.02 for x > 5 mpc is required to explain the T_e([O III]) profile, whereas f_B([S III]) < 0.01 and f_B([N II]) < 0.01 are required at all positions.

It is not surprising that f_B should vary between ions, since the photoionization equilibrium ion fraction of O²⁺ from zone A is much lower than that of N⁺ or S²⁺. Assuming f_B ≪ 1, the ionic abundances given in the "Cut 1, HH 204" column of Table 4 correspond to zone A. These yield O²⁺/O = 0.005 and S²⁺/S = 0.61 if the abundances of unobserved ion stages are negligible. The lack of [N iii] lines means that N⁺/N cannot be estimated directly but is likely of order unity. The fact that O²⁺ is only present in trace amounts in the photoionization equilibrium gas means that the relative contribution from the postshock cooling zone is much larger than for S²⁺ and N⁺. This is confirmed by emission line imaging of HH 204 (Weilbacher et al. 2015), which shows a morphology in [S III] and [N II] that is clearly dominated by the compact jet bullet, whereas the emission in [O III] is more diffuse within the parabolic envelope of the bow shock.

Although part of this high ionization degree gas may be out of photoionization equilibrium, fortunately, its impact is negligible in the global abundance analysis of HH 204. The fact that T_e([N ii]) and T_e([S iii]) are kept in balance in HH 204 proves that the low and medium ionization degree gas, which comprises more than ∼99% of the total, is in photoionization equilibrium.

9.3.2. A Trapped Ionization Front

The detection of emission lines of neutral elements such as [O i] and [N i] and the high density and low degree of ionization of HH 204 suggest that it contains an ionization front. In previous studies, the detection of these lines has been interpreted as a product of the interaction of HH 204 with neutral material, such as that found in Orion's Veil (O'Dell et al. 1997a, 1997b; Takami et al. 2002). However, there are several arguments against this scenario and in favor of the existence of a trapped ionization front. (i) The spatial distribution of the [O i] emission, shown in Figure 2, is more concentrated than that of [O ii] or [O iii], located at the southeast of HH 204, in the opposite direction to θ¹ Ori C, consistent with a zone shielded from the ionizing radiation. (ii) As we discussed in Section 5.3, deuterium lines are produced by fluorescence excitation areas beyond an ionization front. Finally, (iii) the combination of the tangential and radial motions of HH 204 allows one to know the 3D trajectory of its associated jet. From its apparent distance to Orion-S (its likely origin), Doi et al. (2004) estimated that HH 204 has moved ∼0.15 pc radially toward the observer. Although van der Werf et al. (2013) argued that the Orion Veil lies ∼0.3 pc apart from Orion-S, Abel et al. (2016) established that the distance must be significantly larger, and therefore a direct interaction between HH 204 and the Veil is unlikely. If those distance estimations are correct, HH 204 would be located within the main ionized gas volume of the Orion Nebula or interacting with the nearer ionized layer (see Abel et al. 2019; O'Dell et al. 2020).

9.3.3. The ADF and the "True" O Abundance

The origin of the AD problem has been related to temperature, density, or chemical inhomogeneities in the nebulae or fluorescence effects on the intensity of RLs (see Peimbert 1967; Torres-Peimbert et al. 1980; Pequignot et al. 1991; Liu et al. 2001; García-Rojas & Esteban 2007; Escalante et al. 2012, and references therein). As mentioned in Section 5.2, the O⁺ abundances calculated with RLs and CELs are equal in HH 204. Since practically all oxygen is singly ionized in this object, this implies that HH 204, contrary to what is usually found in ionized nebulae, does not show an ADF in total O abundance. Therefore, the "true" O abundance should be ∼8.6 in this object, slightly lower than the recommended solar O abundance (8.73 ± 0.07; Lodders 2019). In this regard, there are three properties of HH 204 that we want to highlight. (i) In Section 5.1, we show that the spatial distribution of T_e([N ii]) is constant; i.e., there are no significant temperature fluctuations in the plane of the sky that may be translated into fluctuations in the line of sight for ions of a low degree of ionization (see Section 9.3.1). The presence of temperature fluctuations would produce the underestimation of the O⁺ abundance based on CELs. (ii) In Section 4.3, we show that the effects of starlight fluorescence are negligible in the determination of the abundances of Ni⁺ and Fe⁺ due to the large distance between HH 204 and the ionizing source in addition to the high density of the HH object. Thus, if there is any mechanism in which the continuum pumping can affect the population of the levels of multiplet 1 of O i, this may be diminished in a similar way. (iii) The jet geometry of HH 204, with a relatively small angle of 32° ± 6° with respect to the plane of the sky (see Section 8), implies that any gradient in the electron density of the gas along the jet axis should be separated in the plane of the sky. This is a more favorable geometry for analysis than the case of a jet flowing directly toward the observer, where different zones will overlap the same line of sight. Therefore, chemical or density inhomogeneities in the line of sight appear not to be present in HH 204. Unfortunately, we cannot perform a similar analysis for the ADF(O²⁺) because O ii RLs are not detected in HH 204.

In Table 8, we compile the O abundances obtained in all chemical abundance studies of the Orion Nebula based on deep echelle spectroscopy taken with UVES. We include determinations based on both RLs and CELs, assuming t² = 0 in the last case. A first note of caution should be given concerning the fraction of O depleted onto dust grains, which may be different in different parts of the nebula. Mesa-Delgado et al. (2009) estimated this fraction to be typically ∼0.12 dex, but it may be lower in HH objects due to destruction of dust grains in shock fronts. This implies a maximum extra uncertainty of ∼0.1 dex in any given O abundance measurement due to depletion variations.

Table 8. Oxygen Abundances in the Orion Nebula Based on UVES Spectroscopy

	RLs			CELs
Region	O+	O2+	O	O+	O2+	O	Reference
Orion Nebula	8.15 ± 0.13	8.57 ± 0.01	8.71 ± 0.03	7.76 ± 0.15	8.43 ± 0.01	8.51 ± 0.03	Esteban et al. (2004)
	8.01 ± 0.12	8.46 ± 0.03	8.59 ± 0.05	8.00 ± 0.06	8.35 ± 0.03	8.51 ± 0.03	Mesa-Delgado et al. (2009)
	8.25 ± 0.06	8.52 ± 0.02	8.70 ± 0.03	7.83 ± 0.05	8.35 ± 0.03	8.46 ± 0.03	Méndez-Delgado et al. (2021)
	⋯	8.40 ± 0.03	8.60 ± 0.03 a	8.18 ± 0.06	8.04 ± 0.02	8.42 ± 0.04	This work
HH 202 S	8.25 ± 0.16	8.44 ± 0.03	8.65 ± 0.05	8.29 ± 0.06	8.08 ± 0.03	8.50 ± 0.04	Mesa-Delgado et al. (2009)
HH 529 II	<7.91	8.83 ± 0.07	8.83 ± 0.07	7.36 ± 0.12	8.54 ± 0.03	8.57 ± 0.03	Méndez-Delgado et al. (2021)
HH 529 III	<7.95	8.84 ± 0.09	8.84 ± 0.09	7.51 ± 0.22	8.48 ± 0.03	8.53 ± 0.03	Méndez-Delgado et al. (2021)
HH 204	8.57 ± 0.03	<7.54	8.57 ± 0.03	8.62 ± 0.05	6.34 ± 0.02	8.62 ± 0.05	This work

Notes. Abundances in units of 12+log(Xⁿ⁺/H⁺) or 12+log(X/H). The bold values are those adopted for the determination of chemical abundances.

^a Using the O⁺ abundance based on CELs.

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If we assume that the O abundances based on RLs are the "true" ones for all objects, then HH 529 II and III show higher O/H ratios than the rest. In Paper I, we discussed the possibility of having a slight overmetallicity in HH 529 II and III due to the entrainment of material from the accretion disk of the stellar source of the jets. On the other hand, the O abundances based on RLs found in the nebular components studied in Paper I and Esteban et al. (2004) are also marginally higher than what is found in HH 204. These discrepancies may have explanations of a different nature in each case, apart from dust depletion variations.

The low measured CEL abundance values found in the more highly ionized regions of the nebula could be reconciled with the HH 204 value by considering different small proportions of O depletion onto dust grains in addition to small contributions from other phenomena, such as temperature fluctuations. In this context, if we assume that half of the difference between ∼8.6—considering that the O abundance obtained in HH 204 is the true one of the Orion Nebula—and the O abundance based on CELs obtained by Esteban et al. (2004) is due to dust depletion and the rest to temperature fluctuations, this would be compatible with t² ∼ 0.008, a value considerable smaller than the t² ∼ 0.022 necessary to match the O abundances from RLs obtained in the same spectrum. In this case, the relevant question is why the RLs are giving higher O abundances in all cases except HH 204. An important difference between the determination of the O abundance in HH 204 and the other zones or objects included in Table 8 is that, in HH 204, the contribution of O²⁺ to the total abundance is negligible. It is important to say that Mesa-Delgado et al. (2009) also obtained an ADF(O⁺) equal to zero in both HH 202 S and the nebular component. However, the contribution of O²⁺ is much larger in those spectra, and their ADF(O²⁺) are not zero. This result suggests that the AD problem is affecting specially to the O²⁺ abundances and perhaps related to unaccounted effects on the intensity of O ii RLs. It is still premature to draw any conclusions in this regard, but we will explore this important issue further in future papers of this series.

From Table 7, it is clear that the Fe abundance in HH 204 is higher than in the other components due to dust destruction at the bow shock. Following the same procedure as in Paper I, comparing the observed Fe/O values in HH 204 and the nebular component with the expected solar value (Lodders 2019), we estimate that ∼6% of the total Fe is in the gaseous phase in the nebular component, while this fraction goes up to 21% in HH 204, representing an increase of a factor 3.5. A similar factor can be assumed for Ni.

9.4. On the Presence of High-density Inclusions

We have studied the physical conditions and chemical composition of the photoionized HH object HH 204 through deep high spectral resolution UVES spectroscopy and HST imaging. Our spectral resolution allows us to cleanly separate HH 204 from the various kinematic components of the Orion Nebula along the same line of sight.

We have analyzed the distribution of the physical conditions of HH 204 along the slit with subarcsecond spatial resolution. We find a steady increase of n_e from ∼10,000 cm⁻³ at ∼13 mpc behind the bow shock to ∼20,000 cm⁻³ close to it. The temperature determined from the most abundant ion stages, such as T_e([N ii]) and T_e([S iii]), is approximately constant at 9000 ± 500 K along the slit. In contrast, T_e([O iii]) is generally higher and shows a pronounced gradient from ≈17,000 K close to the bow shock to ≈12,000 K at distances >5 mpc. We interpret this in terms of a two-zone model (Section 9.3.1). Zone A represents gas that is at the photoionization equilibrium temperature and contributes the overwhelming majority of the low- and intermediate-ionization emission. Zone B is a higher-temperature cooling layer behind the bow shock, and this contributes a significant fraction of the [O iii] emission but contributes negligibly to the other ions.

We estimate that ∼99% of the gas in the observed area of HH 204 is composed of low- and intermediate-ionization stages (ionization potential <25 eV). Based on the intensity of CELs, we determine the ionic abundances of O⁺, O²⁺, N⁺, Ne²⁺, S⁺, S²⁺, Cl⁺, Cl²⁺, Ar²⁺, Fe⁺, Fe²⁺, Ni⁺, Ni²⁺, Ca⁺, and Cr⁺. We also calculate the ionic abundances of He⁺, O⁺, and C²⁺ from the relative intensity of RLs. In HH 204, we find no difference when determining the O⁺ abundance using CELs or RLs. Since practically all O is O⁺ in this object, we can say that the AD is virtually zero for HH 204, contrary to what is found in essentially all ionized nebulae. Both CELs and RLs provide an O abundance of ∼8.60 ± 0.05, slightly lower than the solar value of O = 8.73 ± 0.07 recommended by Lodders (2019) but consistent with many other independent determinations for the Orion Nebula.

Due to the low degree of ionization of HH 204, we can derive the O, N, S, Cl, Fe, and Ni abundances without ICFs. In principle, those O, N, S, and Cl abundances should be representative of the Orion Nebula ones as well. The Fe and Ni abundances of HH 204 are a factor of 3.5 higher than in the Orion Nebula due to the destruction of dust grains at the bow shock. We also found direct evidence of the presence of an ionization front trapped in HH 204, such as the detection of deuterium lines produced by nonionizing far-UV photons.

From archival HST imaging with a higher spatial resolution than our spectra, we find a narrow border of high [O iii]/Hα that traces the leading edge of the bow shock in HH 204 (Figure 18 and Section 5.4). We identify this with the postshock cooling layer, with a width of ≈0.1 mpc. This is the same as the high-temperature zone B, which we invoked in order to explain the spatial profile of T_e([O III]) in the UVES spectra. Note, however, that this layer is much narrower than can be spatially resolved in our spectroscopic observations, which means that the effects on temperature diagnostics are diluted. We predict that much higher values of T_e([O III]) ≈ 30,000 K would be found if the λ4363/λ5007 ratio were to be observed at a spatial resolution of 005.

We investigate the origin of the driving jets of both HH 204 and the nearby HH 203 using both proper motions and channel maps of highly blueshifted emission (Section 8). We find that HH 203 is the superposition of two flows: a high-ionization and high-velocity flow at PA 118, which originates in the Orion-S region, plus a low-ionization and lower-velocity flow at PA 140, which originates near the Trapezium. The proper motion of the HH 204 bow shock is closely aligned with the PA 140 flow, suggesting that HH 204 may also be driven by this same jet, but there is little evidence of such a connection from the blueshifted channel maps. Instead, there is evidence for a third flow at PA 108 that appears to be feeding into HH 204 and may be connected to a molecular hydrogen filament originating in the region to the south of Orion-S.

Our observations allow us to separate and analyze the spectrum of the DBL, an ionized gas component with a radial velocity different from that of the Orion Nebula and HH 204. We have estimated its physical conditions, its T_e, for the first time, revealing that it has a density lower than the Orion Nebula. We have calculated its chemical composition for the first time.

Our analysis of the spectrum of the kinematic component corresponding to the Orion Nebula reveals a lower ionization degree and n_e with respect to the results of Paper I. This comparison also indicates that T_e in the Orion Nebula decreases with the radial distance from θ¹ Ori C. The chemical composition of the nebular component is similar to that found in Paper I, although there seems to be a slightly lower O abundance (less than 0.04 dex), perhaps related to different depletion factors onto dust grains of this element.

We carry out the exercise of simulating a spectrum with a lower spectral and spatial resolution, where the spectra of the different kinematic components are mixed. We find that the analysis of this integrated spectrum can lead to erroneous physical conditions and chemical abundances. For example, the estimation of n_e by averaging n_e([O ii]), n_e([S ii]), and n_e([Cl iii]) underestimates the true density, resulting in an overestimation of the temperature of the low-ionization ions, which constitute an important fraction of the gas in HH 204, the dominant component of the integrated spectrum. This fact leads to an underestimate of the abundances and to obtaining a mistaken average degree of ionization, the parameter on which most ICF schemes are based. Therefore, the determination of the chemical abundances would be wrong in practically all elements. Indicators of density such as I([Fe III] λ4658)/I([Fe III] λ4702) may be used to detect the presence of high-density clumps associated with HH objects or shocks. A similar point is made by O'Dell et al. (2021) with respect to unrecognized heterogeneity in physical conditions leading to misleading results, and we echo the warning of that paper.

This work is based on observations collected at the European Southern Observatory, Chile, proposal No. ESO 092.C-0323(A). We are grateful to the anonymous referee for helpful comments. We acknowledge support from the Agencia Estatal de Investigación del Ministerio de Ciencia e Innovación (AEI-MCINN) under grant Espectroscopía de campo integral de regiones H ii locales. Modelos para el estudio de regiones H ii extragalácticas with reference 10.13039/501100011033. W.J.H. is grateful for financial support provided by Dirección General de Asuntos del Personal Académico, Universidad Nacional Autónoma de México, through grant Programa de Apoyo a Proyectos de Investigación e Inovación Tecnológica IN107019. J.G.-R. acknowledges support from Advanced Fellowships under the Severo Ochoa excellence programs SEV-2015-0548 and CEX2019-000920-S. J.E.M.-D. is grateful for the support of the Instituto de Astrofísica de Canarias under the Astrophysicist Resident Program and acknowledges support from the Mexican CONACyT (grant CVU 602402). The authors acknowledge support under grant P/308614 financed by funds transferred from the Spanish Ministry of Science, Innovation and Universities, charged to the General State Budgets, and with funds transferred from the General Budgets of the Autonomous Community of the Canary Islands by the MCIU.

Generally, the discrepancy between the physical conditions derived from diagnostics based on [Fe iii] lines and those estimated from other ions has been interpreted as a result of errors in the transition probabilities and/or the collision strengths of the Fe²⁺ ion (see the introduction of Laha et al. 2017). With HH 204, we have an excellent opportunity to test the reliability of the atomic data we use for this ion for the following reasons. (i) We have enough spectral resolution to separate the emission of HH 204 from that of the Orion Nebula. (ii) Owing to its geometry and 3D trajectory, we do not expect significant inhomogeneities in the physical conditions within the line of sight for ions of low and intermediate ionization stages. Due to this, all of the density diagnostics used in Table 3 are consistent with each other, while the global gas temperature, due to the low degree of ionization, is well represented by T_e([N ii]). (iii) The [Fe iii] emission is enhanced owing to the destruction of dust grains containing Fe atoms in the shock, which allows one to have a good signal-to-noise ratio even for some weak lines that are difficult to detect.

We have used a set of transition probabilities compiled in PyNeb, which includes the data from Quinet (1996) and Johansson et al. (2000) for ⁵D − ⁵S₂ transitions. However, these transitions produce lines out of the spectral range covered by our observation, so we finally only use the calculations from Quinet (1996). Table 12 shows that the transition probabilities we use are in good agreement with the observed intensity ratios of lines arising from the same upper level in the case of lines used to determine physical conditions. However, the intensity ratios between lines that arise from different upper levels do depend on n_e and T_e. As discussed above and in Section 9.4, in HH 204, there are no significant density inhomogeneities that may produce a bias in some diagnostics, contrary to the case analyzed in Section 7. Thus, all of the density diagnostics included in Table 3 give consistent results, and the average of n_e([O ii]), n_e([S ii]), n_e([Cl iii]), and n_e([Fe ii]) is n_e = 13,330 ± 550. Using this density for its calculation, T_e([N ii]) remains practically unchanged from what is shown in Table 3. Considering these values of n_e and T_e([N ii]), we can check the validity of [Fe iii] atomic data by applying the procedure that we describe below. First, we take into account all of the observed [Fe iii] lines that are not affected by blending with other lines, sky features, or telluric absorptions. Then we normalize their emission with respect to I([Fe III] λ4658)/I(Hβ) = 1000. We discard the [Fe iii] λ λ3355.50, 7078.22 lines, since their FWHMs are much wider than the rest of the [Fe iii] lines, which is indicative of line blending. The [Fe iii] λ9203.85 line is also discarded because it shows a radial velocity of ∼10 km s⁻¹, larger than the velocities of the rest of the [Fe iii] lines, which may be indicative of a doubtful identification. We also discard [Fe iii] λ8838.14 because, although we deblend it from a very close sky feature, its intensity may not be completely reliable. Once we have the set of [Fe iii] lines with confident observed intensity ratios, they are compared with the predictions of the atomic data for the assumed physical conditions and considering error propagation. The results are shown in Table 14.

Table 14 does not include ⁵D − ⁷S transitions ([Fe iii] λ λ3322.47, 3371.35, 3406.18 lines) because their transition probabilities are not calculated in the reference of the atomic data used (their "Predicted" and "Difference" columns are empty). However, their measured intensities can be used to check other atomic data sets that do include them. In general, Table 14 shows good agreement between the predicted and observed intensity ratios of [Fe iii] lines. Only four lines (λ λ4008.34, 4079.69, 4985.88, and 7088.46) show differences larger than 10%, exceeding the error bars. This can be attributed to errors in their atomic data. The first two lines arise from the same ³G₄ upper level, so their intensity ratio only depends on their transition probabilities. Although the I([Fe III] λ4008.34)/I([Fe III] λ4079.69) ratio is not included in Table 12—these lines were not used to determine physical conditions—its intensity ratio of 4.43 ± 0.30 is larger than the theoretical one of 3.92. Therefore, it is plausible that part of the observed discrepancy is due to incorrect transition probabilities. The largest differences reported in Table 14 are for the [Fe iii] λ λ4985.88, 7088.46 lines, but we cannot find an obvious explanation for this. In addition to the atomic data used, we have checked other sets: for transition probabilities, Nahar & Pradhan (1996) and Bautista et al. (2010), and for collision strengths, Bautista et al. (2010) and Badnell & Ballance (2014). We have tried all possible combinations of these data. Of the nine combinations, the atomic data we use in this paper minimize the difference between the predicted and measured intensity ratios. The results of this Appendix indicate that the atomic data used in this work for [Fe iii] lines contribute little to errors in the derived physical conditions and Fe²⁺ abundances, at least for the conditions of HH 204. As we discuss in Section 7, the discrepancy normally found between n_e([Fe iii]) and the classical diagnostics—such as n_e([O II]) or n_e([S II])—may be rather indicative of the presence of high-density inclusions within the line of sight. For a complete test of the atomic data, similar studies would be necessary in different ranges of physical conditions. We will continue investigating this topic in other HH objects in future papers of this series.

In this Appendix, we include the following material.

• 1.  
  Figure 25. Plasma diagnostics for the individual components analyzed in this work.
• 2.  
  Table 9. Sample of lines of the spectra of cut 1. The full version is included in the .tar.gz package in machine-readable format. The package also contains the cut 2, combined, and individual blue and red pixel data.
• 3.  
  Table 10. Atomic data set used for CELs.
• 4.  
  Table 11. Atomic data set used for RLs.
• 5.  
  Table 12. Measured and predicted [Fe iii] intensity ratios from lines that arise from a common upper level.
• 6.  
  Table 13. Measured and predicted [Fe ii] intensity ratios from lines that arise from a common upper level.
• 7.  
  Table 14. Measured and predicted [Fe iii] intensity ratios for all detected lines using the atomic data chosen in this work.
• 8.  
  Table 15. Pixel-to-pixel spatial distribution of the physical conditions and ionic abundances of HH 204 in the UVES blue arm spectra.
• 9.  
  Table 16. Pixel-to-pixel spatial distribution of the physical conditions and ionic abundances of HH 204 in the UVES red arm spectra.
• 10.  
  Table 17. Pixel-to-pixel spatial distribution of ionic abundances of HH 204 in the UVES red arm spectra.

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Table 9. Sample of 15 Lines from the Spectra of Cut 1

		HH 204						The Orion Nebula+The DBL
λ0 (Å)	Ion	λobs	Vel. $\left({\lambda }_{0}\right)$ (km s−1)	FWHM (km s−1)	$F\left(\lambda \right)$/$F\left({\rm{H}}\beta \right)$	$I\left(\lambda \right)$/$I\left({\rm{H}}\beta \right)$	Err. %	λobs	Vel. $\left({\lambda }_{0}\right)$ (km s−1)	FWHM (km s−1)	$F\left(\lambda \right)$/$F\left({\rm{H}}\beta \right)$	$I\left(\lambda \right)$/$I\left({\rm{H}}\beta \right)$	Err. %	Notes
4701.64	[Fe III]	4701.25	−24.76	18.30 ± 0.02	1.424	1.460	2	4701.83	12.22	23.27 ± 0.33	0.193	0.197	3
4713.14	He I	4712.80	−21.48	18.57 ± 0.31	0.206	0.211	3	4713.35	13.51	27.03 ± 0.30	0.436	0.443	3
4728.07	[Fe II]	4727.75	−20.08	13.57 ± 0.36	0.072	0.073	4	4728.45	24.31	8.18 ± 4.17	0.004	0.004	30
4734.00	[Fe III]	4733.57	−27.00	18.18 ± 0.05	0.634	0.647	2	4734.15	9.74	23.81 ± 1.68	0.067	0.068	6
4740.17	[Ar IV]	*	*	*	*	*	*	4740.35	11.65	14.86 ± 6.68	0.006	0.006	29
4754.81	[Fe III]	4754.42	−24.90	18.29 ± 0.04	0.800	0.813	2	4755.00	11.67	22.95 ± 0.58	0.127	0.128	4
4769.53	[Fe III]	4769.14	−24.77	18.23 ± 0.04	0.505	0.512	2	4769.73	12.32	22.75 ± 0.72	0.064	0.065	4
4774.73	[Fe II]	4774.42	−19.69	13.19 ± 0.51	0.070	0.071	4	*	*	*	*	*	*
4777.70	[Fe III]	4777.38	−20.30	17.82 ± 0.21	0.304	0.308	3	4777.97	16.73	24.34 ± 2.76	0.032	0.032	10
4803.29	N II	*	*	*	*	*	*	4803.46	10.50	18.10 ± 2.08	0.019	0.019	8
4814.54	[Fe II]	4814.23	−19.37	14.38 ± 0.04	0.393	0.396	2	*	*	*	*	*	*
4874.50	[Fe II]	4874.18	−19.51	13.84 ± 0.54	0.051	0.051	4	*	*	*	*	*	*
4861.32	H I	4860.97	−21.47	24.67 ± 0.01	100.000	100.000	2	4861.52	12.45	30.59 ± 0.01	100.000	100.000	2
4861.32	H I	4859.66	−102.26	14.31 ± 1.90	0.106	0.106	10	*	*	*	*	*	*	Deuterium
4874.50	[Fe II]	4874.18	−19.51	13.84 ± 0.54	0.051	0.051	4	*	*	*	*	*	*
4881.07	[Fe III]	4880.71	−21.92	18.00 ± 0.01	2.251	2.245	2	4881.30	14.32	20.70 ± 0.17	0.248	0.247	3

Note. The full version is included in the .tar.gz package, along with the cut 2, combined, and individual blue and red pixel data.

Only a portion of this table is shown here to demonstrate its form and content. A machine-readable version of the full table is available.

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