$\chi^2$ fit from the LTE modeling process. As Table 6 shows, we modeled the H$_2$S emission from the ortho and para spin modifications separately. In both cases, the narrow component has a line width of ~3.6 km s$^{-1}$, a low excitation temperature (24 K), and is fairly extended (emission extending beyond the Herschel beam) whereas the broad component has a larger line width (~8 km s$^{-1}$), is warmer (~80 K), and yet is still fairly extended (>35"). The ortho to para ratio in the narrow component is 1.1 ± 0.3 and in the broad component is 0.9 ± 0.1 indicating a spin temperature of 9 ± 2 K. These values are consistent with those determined for H$_2$O in Orion-S by Choi et al. (2014) and for Formaldehyde (above). Similarly, this low ortho to para ratio suggests either a very low formation temperature for H$_2$S or that non-LTE effects had an important role in its formation. To model H$_2$$^{34}$S, we coupled its single ortho transition with those of the common isotopologue. The isotopic ratio is another possible free parameter in CASSIS, which assumes that the other parameters of both isotopologues are identical. The isotopic ratio converged to 31 ± 9. Note, however, that there is only one weak transition of H$_2$$^{34}$S that we used to determine this ratio. For comparison, typical values for the $^{32}$S/$^{34}$S ratio in galactic molecular clouds are ~19 ± 8 (Lucas & Liszt 1998), which is consistent with the solar value of 23 (Anders & Grevesse 1989).

**CS:** CS is a well-known tracer of dense gas due to its high critical density (Plume et al. 1992). We observe many transitions of CS, from *J* = 10–9 to *J* = 19–18. Like CO, CS requires two components to be successfully fit by a Gaussian profile and modeled; one being narrow ($\Delta V_{\text{FWHM}} \sim 4.1$ km s$^{-1}$), extended ($\Omega = 67''$), and cool ($T = 37$ K), and the other being broader ($\Delta V_{\text{FWHM}} \sim 10$ km s$^{-1}$), moderately extended ($\Omega = 35''$), and warm ($T = 108$ K). The broad component is not seen in the spectra of the CS isotopologues but was needed to successfully model the transitions. $^{13}$CS and C$^{34}$S were modeled in a fashion similar to that of H$_2$$^{34}$S (i.e., the transitions of the isotopologues modeled simultaneously with those of the common isotope, and found to have an isotopic ratio of 46 ± 17 and 14 ± 5 respectively). The Gaussian fitted (Table 5) and modeled velocity (Table 6) of both components is ~7 km s$^{-1}$, suggesting that both components originate in the same material (Figure 14).

**HCN:** the transitions of HCN (*J* = 6–5 to 13–12) exhibit the characteristic broad-line wing that required both two component Gaussian fitting and LTE modeling (see Table 6). As is usual for other species, the broad component is hotter (67 K versus 34 K), broader (13.4 km s$^{-1}$ versus 4.4 km s$^{-1}$), and less spatially extended (41" versus 64") than the narrow component. While the broad component's $V_{\text{LSR}}$ is lower than that of narrow component (6.6 km s$^{-1}$ versus 7.2 km s$^{-1}$) both are consistent with the systemic velocity of Orion-S (Figure 14).

**HCO$^+$:** HCO$^+$ is another well-known tracer of both dense molecular gas and outflows. We detected the *J* = 6–5 to *J* = 13–12 transitions of HCO$^+$ in Orion-S, which span a wide range of physical conditions ($E_{\text{up}} = 90$ K, $n_{\text{crit}} \sim 3.2 \times 10^7$ cm$^{-3}$ to $E_{\text{up}} = 389$ K, $n_{\text{crit}} \sim 6.0 \times 10^8$ cm$^{-3}$). HCO$^+$ has some of the strongest lines seen in our survey of Orion-S. Despite the fact that we fit the HCO$^+$ transitions with a single Gaussian in Table 4, HCO$^+$ required a second physical component in order to obtain a good $\chi^2$ fit from the LTE modeling process (see Table 6). Both components have a

velocity of ~7 km s$^{-1}$ and are fairly warm (~69 K), suggesting a common origin.

**NH$_3$:** we detect two transitions above the 5$\sigma$ level in Orion-S: one in emission (10,0 – 0,1) and one in absorption (2,1 – 1,0). In fact, we detect three additional transitions of NH$_3$ (all in absorption), but since they were just below the 5$\sigma$ level, we do not report them in Table 4. Modeling these transitions simultaneously requires two components: a cold (*T* ~ 20 K), quiescent ($\Delta V \sim 4.2$ km s$^{-1}$) layer of gas in front of a warmer (*T* ~ 36 K), broader ($\Delta V = 10$ km s$^{-1}$) component. The $V_{\text{LSR}}$ values of both NH$_3$ components are consistent with the systemic velocity of the cloud (Figure 14). The presence of absorption lines provides additional evidence for the existence of two components in Orion-S (one warm and one cool). The fact that the low energy transition is seen in emission, whereas the higher energy transition is seen in absorption may be related to the beam size and the strength of the continuum. At 572 GHz, the continuum is weaker than it is at 1215 GHz (see Figure 13) and may be too beam-diluted to see absorption. However, at higher frequencies, the beam couples better to the source and absorption may become more prevalent. This suggests that NH$_3$ may not come from a high density, hot region, which is consistent with our conclusion that there are no hot cores in Orion-S.

**SO$_2$:** SO$_2$ is also often a tracer of shocks, since it can freeze onto grain mantles at early evolutionary times when the gas is cold and dense, and later be returned to the gas phase by shocks (e.g., Millar & Herbst 1990; Esplugues et al. 2013). Toward Orion-S, we detected a number of SO$_2$ transitions above the 5$\sigma$ level (see Table 4). Given the fact that all the observed SO$_2$ lines are relatively weak, one-component modeling matches the observations remarkably well. LTE modeling for this species (Table 5) shows that the SO$_2$ is fairly warm ($T_{\text{ex}} \sim 150$ K), broad ($\Delta V \sim 6.7$ km s$^{-1}$), and extended (~65"). The velocity of SO$_2$ is also similar to the velocity of the quiescent gas (Figure 14).

**SO:** in contrast with SO$_2$, and despite the fact that we fit the SO transitions with a single Gaussian in Table 4, SO required a second physical component in order to obtain a good $\chi^2$ fit from the LTE modeling process (see Table 6). The narrow component has a line width of 3.8 km s$^{-1}$, a moderately low temperature (34 K), and is extended beyond the Herschel beam. Despite the fact that the broad component ($\Delta V_{\text{FWHM}} = 11$ km s$^{-1}$) is much warmer (122 K), it is still quite extended in size (35").

The large line widths (7–13 km s$^{-1}$), high temperatures (70–150 K), and extended size (>30") determined for the second physical components of these species suggest that the embedded outflows seen in Orion-S (Schmid-Burgk et al. 1990; Ziurys et al. 1990; Zapata et al. 2005) have affected a large volume of the region both thermally and dynamically. Whether or not these shocks have affected the chemistry of the gas will be examined in Section 3.3.

### 3.3. Chemical Comparison with Orion-KL

One of the main goals of this project is to explore the chemical differences and similarities between Orion-S and Orion-KL. As mentioned in Section 1, a detailed comparison between the chemical abundances in Orion-S and Orion-KL is useful, since both regions presumably formed under similar conditions, but could have developed very different chemical abundances based on differences in their ages, densities,
<table><caption>**Table 7**
The Abundance of Species with Respect to H₂</caption><thead><tr><th>Species</th><th>Narrow Comp wrt H<sub>2</sub></th><th>Broad Comp wrt H<sub>2</sub></th><th>Abundance Enhancement Factor</th></tr></thead><tbody><tr><td>CCH</td><td>4.32 × 10<sup>-9</sup></td><td>...</td><td>...</td></tr><tr><td>CN</td><td>1.26 × 10<sup>-9</sup></td><td>...</td><td>...</td></tr><tr><td>C<sup>17</sup>O</td><td>6.80 × 10<sup>-8</sup></td><td>...</td><td>...</td></tr><tr><td><sup>13</sup>C<sup>18</sup>O</td><td>3.93 × 10<sup>-9</sup></td><td>...</td><td>...</td></tr><tr><td>DCN</td><td>2.38 × 10<sup>-11</sup></td><td>...</td><td>...</td></tr><tr><td>HCl</td><td>7.77 × 10<sup>-10</sup></td><td>...</td><td>...</td></tr><tr><td>H<sup>37</sup>Cl</td><td>3.25 × 10<sup>-10</sup></td><td>...</td><td>...</td></tr><tr><td>HNC</td><td>1.17 × 10<sup>-10</sup></td><td>...</td><td>...</td></tr><tr><td>NO</td><td>3.40 × 10<sup>-8</sup></td><td>...</td><td>...</td></tr><tr><td>H<sup>13</sup>CO<sup>+</sup></td><td>3.45 × 10<sup>-11</sup></td><td>...</td><td>...</td></tr><tr><td>H<sub>2</sub>CS</td><td>3.01 × 10<sup>-10</sup></td><td>...</td><td>...</td></tr><tr><td>SO<sub>2</sub></td><td>7.29 × 10<sup>-10</sup></td><td>...</td><td>...</td></tr><tr><td>H<sub>2</sub>S</td><td>9.42 × 10<sup>-9</sup></td><td>3.29 × 10<sup>-7</sup></td><td>35</td></tr><tr><td>H<sup>34</sup>S</td><td>1.46 × 10<sup>-10</sup></td><td>3.41 × 10<sup>-9</sup></td><td>23</td></tr><tr><td>SO</td><td>4.86 × 10<sup>-9</sup></td><td>2.47 × 10<sup>-7</sup></td><td>51</td></tr><tr><td><sup>13</sup>CS</td><td>6.80 × 10<sup>-11</sup></td><td>1.88 × 10<sup>-9</sup></td><td>28</td></tr><tr><td>C<sup>34</sup>S</td><td>2.19 × 10<sup>-10</sup></td><td>4.59 × 10<sup>-9</sup></td><td>21</td></tr><tr><td>NH<sub>3</sub></td><td>6.80 × 10<sup>-10</sup></td><td>9.88 × 10<sup>-8</sup></td><td>145</td></tr><tr><td>H<sub>2</sub>CO</td><td>1.70 × 10<sup>-9</sup></td><td>1.24 × 10<sup>-7</sup></td><td>73</td></tr></tbody></table>

<figure><img src="image_14.png" /><figcaption>**Figure 15.** Abundance (with respect to H₂) of species listed in Tables 5 and 6. Open red triangles indicate the abundance ratio for species fitted by a single component LTE model in Orion-S. In cases requiring two component LTE models for the Orion-S data, the abundance ratio for the narrow component is indicated by solid red triangles and that for the broad component is indicated by solid green squares. Therefore, the dotted line connects species/components that likely trace quiescent gas, whereas the dashed line connects species/components that may trace shocked gas.</figcaption></figure>

temperatures, radiation fields, etc. As part of the HEXOS survey, and in a direct analogue to our study, Crockett et al. (2014) have observed the same frequency range in Orion-KL using the same instrument. Therefore, we have a perfectly matched database with which to compare our results. In this section, we will compare the chemistries of these two regions. All of the Orion-KL data are taken from the HEXOS survey of this region as listed in Crockett et al. (2014) using the column densities derived from their XCLASS LTE modeling of the data.

### 3.3.1. Chemical Abundances in Orion-S

Our LTE modeling produces column densities but, to obtain chemical abundances, we must scale each column density by

<figure><img src="image_15.png" /><figcaption>**Figure 16.** Comparison of the abundances of species detected in the Orion-KL Hot Core to those in Orion-S as given by Equation (1) in Section 3.3.2. Symbols are the same as described in Figure 15.</figcaption></figure>

the H₂ column density, which is not known directly. Therefore, we use C<sup>18</sup>O as a proxy for the H₂ column density. The C<sup>18</sup>O transitions we observe in Orion-S are optically thin and have only a single narrow component (Table 4) that is well fit by a column density of 3.5 × 10<sup>16</sup> cm<sup>-2</sup> (Table 5). To convert this to an H₂ column density, we use a C<sup>18</sup>O:H₂ conversion factor of 1.7 × 10<sup>-7</sup> (Goldsmith et al. 1997) to obtain N<sub>H₂</sub> = 2.1 × 10<sup>23</sup> cm<sup>-2</sup>.

Other species have been modeled using both a narrow and broad component (Table 6), the latter is possibly indicative of gas affected by shock. To compute the chemical abundances in the broad component of species listed in Table 6, we need an estimate of the C<sup>18</sup>O column density specifically present in this broad component. Since the broad component is too optically thin to be detected in C<sup>18</sup>O or even <sup>13</sup>CO, we rely on <sup>12</sup>CO instead for this purpose. In this case, we attempt to model the broad (i.e., line wing) component of our observed <sup>12</sup>CO lines by fitting three Gaussian components to each transition. The first two Gaussians fit the “main” component of the asymmetric CO profile (see Table 4) and the third fits the outflow (ΔV ~ 18 km s<sup>-1</sup>). This third component is then modeled via our LTE procedure to estimate the physical parameters of the <sup>12</sup>CO outflow (T<sub>ex</sub> ~ 200 K; N(CO) ~ 7.1 × 10<sup>16</sup> cm<sup>-2</sup>; Ω > 30′). Dividing this column density by the canonical <sup>12</sup>CO:C<sup>18</sup>O abundance ratio of 500:1 provides a C<sup>18</sup>O column density of 1.4 × 10<sup>14</sup> cm<sup>-2</sup> for the broad component in Orion-S, which is a factor of ~250 smaller than the C<sup>18</sup>O column density in the narrow component as measured by directly modeling our C<sup>18</sup>O observations. Using the same C<sup>18</sup>O:H₂ scaling relationship as above, we obtain N<sub>H₂</sub> = 8.2 × 10<sup>20</sup> cm<sup>-2</sup>.

Therefore, dividing the modeled column densities of the broad components of the species in Table 6 by 8.2 × 10<sup>20</sup> cm<sup>-2</sup> gives the abundance (with respect to H₂) of all species in Orion-S that are possibly affected by shocks. Dividing the modeled column densities of the rest of the species in Table 5, as well as the narrow component of the species in Table 6, by 2.1 × 10<sup>23</sup> cm<sup>-2</sup> gives the abundance (with respect to H₂) of all species in Orion-S that likely originate in quiescent gas. The results are provided in Table 7 and illustrated in Figure 15, which clearly show that in the broad component (green squares) the abundances are enhanced
by a factor of 10–100 with respect to the narrow component (red triangles). Abundance enhancements of this magnitude indicate classic shock behavior (Bachiller & Pérez Gutiérrez 1997). This suggests that shock chemistry is playing an important role in Orion-S.

### 3.3.2. Species Common to Both Orion-KL and Orion-S

Crockett et al. (2014) detected ~13,000 lines from 39 different molecules (79 species if one includes all the isotopologues). This is considerably more than the 685 lines from 52 species (including isotopologues) that we have detected in Orion-S. In addition, the lines in Orion-KL are typically an order of magnitude stronger than those seen in Orion-S. A more interesting comparison, however, is to examine the abundances of species common to the two sources.

Crockett et al. (2014) produce column densities for many of the detected species in Orion-KL. To compare with the chemical abundances in Orion-S (Section 3.3.1), we must also scale by the H₂ column density in Orion-KL. For Orion-KL, we use the C¹⁸O column densities derived by Plume et al. (2012) as a proxy, which breaks down the results for each of the four known kinematic components: the Hot Core ($V_{LSR} \sim 4-6$ and $\Delta V \sim 7-12$ km s⁻¹), the Plateau ($V_{LSR} \sim 7-11$ and $\Delta V_{LSR} \ge 20$ km s⁻¹), the Compact Ridge ($V_{LSR} \sim 7-9$ and $\Delta V \sim 3-6$ km s⁻¹), and the Extended Ridge ($V_{LSR} \sim 8-10$ and $\Delta V \sim 2-4$ km s⁻¹) (Blake et al. 1987). The H₂ column density can then be produced using the same C¹⁸O:H₂ conversion factor of $1.7 \times 10^{-7}$.

Producing abundances in this way does depend on the assumptions regarding the C¹⁸O:H₂ abundance ratio. However, by dividing the abundance of a given species in Orion-KL by the abundance of the same species in Orion-S, we eliminate the C¹⁸O:H₂ abundance ratio altogether and are essentially normalizing to the C¹⁸O column density in each region. This does, of course, assume that C¹⁸O abundances are the same in both sources, which may be reasonable based upon similarities between the observed C¹⁸O:C¹⁷O ratios (e.g., 2.5 in Orion-S, 3.0 in the Hot Core, 6.5 in the Compact Ridge, 3.5 in the Plateau, and 2.3 in the Extended Ridge). These ratios are also consistent with those found by Ladd (2004).

Therefore, we are essentially creating the following ratio.

$$ \frac{X_{KL}}{X_S} = \frac{\left( \frac{N_i}{N_{C^{18}O}} \right)_{KL}}{\left( \frac{N_i}{N_{C^{18}O}} \right)_S} \quad (1) $$

where $N_i$ is the column density of species i and $N_{C^{18}O}$ is the column density of C¹⁸O. The subscripts KL and S refer to this ratio in Orion-KL and Orion-S respectively. Given the four distinct kinematic components of Orion-KL, we create this ratio for the Hot Core (HC), the Plateau (P), the Compact Ridge (CR), and the Extended Ridge (ER) separately and, again, use different values for the C¹⁸O abundance in Orion-S depending on whether the species in question has a narrow or broad spectral line profile.

Figure 16 shows the comparison between Orion-S and the Orion-KL Hot Core. Note that, by common use, the term “hot core” refers to a dense, warm region surrounding a central high-mass protostellar object that dominates its energetics (Kurtz et al. 2000, p. 299). It has been argued that the eponymous hot core in the Orion-KL region does not fulfill this

<figure><img src="image_16.png" /><figcaption>Figure 17. Same as for Figure 16 except for the Orion-KL Compact Ridge.</figcaption></figure>

<figure><img src="image_17.png" /><figcaption>Figure 18. Same as for Figure 16 except for the Orion-KL Plateau.</figcaption></figure>

<figure><img src="image_18.png" /><figcaption>Figure 19. Same as for Figure 16 except for the Orion-KL Extended Ridge.</figcaption></figure>

criterion (Zapata et al. 2011). Rather, these authors suggest that this region is instead powered by the aftermath of the explosion caused by a stellar merger event (Bally & Zinnecker 2005). Regardless of this, in the present paper, we are comparing the chemical abundances of Orion-S with those in what is
traditionally referred to as the “hot-core” component of Orion KL. The x-axis indicates the species and the y-axis shows the ratio as calculated from Equation (1). Open red triangles indicate molecules for which one-component LTE models in Orion-S were sufficient. In cases where we required two components to model the Orion-S data, the solid red triangles indicate the ratio for the narrow component and the solid green squares indicate the ratio for the broad component. The dotted line connects species/components that likely trace quiescent gas, whereas the dashed line connects species that have a broad component. Not all species detected and modeled in Orion-S are represented in this figure. This is due to the fact that Crockett et al. (2014) did not model all their detected species (notably the atomic species), nor did they provide column densities for species in which the lines were optically thick (e.g., CO, $^{13}$CO, HCO$^{+}$, CS) in Orion-KL. Note also that not every species listed in Figure 16 has a symbol associated with it (e.g., CN, HCl, SO, etc.). This is because emissions from these species were not attributed to the HC, but rather to one or more of the other kinematic components in Orion-KL. Error bars are calculated from the statistical uncertainties determined from our LTE modeling of Orion-S (~10%–50%; see Tables 5 and 6) and with the assumption of 10% error bars of the reported column densities in Orion-KL Crockett et al. (2014), which includes the effects of calibration errors, pointing errors, etc. However, to account for the possibility that LTE is not a good approximation in either Orion-S or Orion-KL, we add an additional 40% error to the column densities. This value is based on a comparison of LTE versus non-LTE column density calculations for the Orion-KL Extended Ridge (Crockett et al. 2014).

An inspection of Figure 16 clearly shows that the abundances of species in the Orion-KL HC are significantly higher than those in the narrow component of Orion-S (dotted line in Figure 16). Except for CCH, C$^{17}$O, and H$^{13}$CO$^{+}$, the abundances in the HC are ≳10 times larger than those in Orion-S. Examining the abundance ratios in the narrow component, we obtain $\langle \frac{X_{KL}}{X_S} \rangle = 135$ (SD = 260), where SD is the standard deviation about the mean. The large standard deviation simply reflects the enormous scatter in the ratios (note that the y-axis in Figure 16 is on the log scale). Although still not a good match, the disagreement is smaller for species that have a broad component, (dashed line in Figure 16). In this case, we obtain $\langle \frac{X_{KL}}{X_S} \rangle = 6$ (SD = 12). Given the lack of complex molecules noted in Section 3.2.2 and the poor match in the abundances between the Orion-KL HC and Orion-S, this suggests that the gas detected in this study of Orion-S does not originate in a hot core.

Figure 17 shows that the abundances of species in the Orion-KL CR are also higher than those in the narrow component of Orion-S ($\langle \frac{X_{KL}}{X_S} \rangle = 23$; SD = 45) but the agreement is better than it is for the HC. Again, the match is better to the broad component ($\langle \frac{X_{KL}}{X_S} \rangle = 1$; SD = 2) of the two component fits in Orion-S (dashed line) than it is to the narrow component.

Figure 18 shows the comparison between Orion-S and the Orion-KL Plateau region. The match between abundances here is clearly better than it is for the HC or the CR with $\langle \frac{X_{KL}}{X_S} \rangle = 14$ (SD = 22) for the narrow component (dotted line) and 1 (SD = 2) for the broad component (dashed line). Note that, although

<table><caption>**Table 8**

Column Density Upper Limit for Species in Orion-KL Not Detected in Orion-S</caption><thead><tr><th>Species</th><th>N<sub>t</sub>, T<sub>ex</sub>=40 K</th><th>N<sub>t</sub>, T<sub>ex</sub>=80 K</th><th>E<sub>u</sub></th></tr></thead><tbody><tr><td><sup>15</sup>NH<sub>3</sub></td><td>7.0 × 10<sup>11</sup></td><td>3.0 × 10<sup>12</sup></td><td>700</td></tr><tr><td><sup>29</sup>SiO</td><td>2.0 × 10<sup>12</sup></td><td>1.0 × 10<sup>12</sup></td><td>700</td></tr><tr><td><sup>30</sup>SiO</td><td>3.0 × 10<sup>12</sup></td><td>2.0 × 10<sup>12</sup></td><td>300</td></tr><tr><td>C<sub>2</sub>H<sub>3</sub>CN</td><td>N/A</td><td>N/A</td><td>300</td></tr><tr><td>C<sub>2</sub>H<sub>5</sub>CN</td><td>N/A</td><td>N/A</td><td>...</td></tr><tr><td>C<sub>2</sub>H<sub>5</sub>OH</td><td>N/A</td><td>N/A</td><td>...</td></tr><tr><td>C<sup>33</sup>S</td><td>6.0 × 10<sup>12</sup></td><td>4.0 × 10<sup>12</sup></td><td>700</td></tr><tr><td>CH<sub>2</sub>DOH</td><td>6.0 × 10<sup>13</sup></td><td>6.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td>CH<sub>3</sub>NH</td><td>7.0 × 10<sup>12</sup></td><td>2.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td>CH<sub>3</sub><sup>13</sup>CN</td><td>3.0 × 10<sup>14</sup></td><td>2.0 × 10<sup>13</sup></td><td>700</td></tr><tr><td>CH<sub>3</sub>CN</td><td>7.0 × 10<sup>14</sup></td><td>5.0 × 10<sup>13</sup></td><td>700</td></tr><tr><td>CH<sub>3</sub>CN, v<sub>8</sub> = 1</td><td>N/A</td><td></td><td>700</td></tr><tr><td><sup>13</sup>CH<sub>3</sub>CN</td><td>4.0 × 10<sup>14</sup></td><td>2.0 × 10<sup>13</sup></td><td>700</td></tr><tr><td>CH<sub>3</sub>OCHO</td><td>2.0 × 10<sup>15</sup></td><td>2.0 × 10<sup>15</sup></td><td>300</td></tr><tr><td>CH<sub>3</sub>OD</td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td><sup>13</sup>CH<sub>3</sub>OH</td><td>N/A</td><td>N/A</td><td>300</td></tr><tr><td>D<sub>2</sub>O</td><td>5.0 × 10<sup>11</sup></td><td>2.0 × 10<sup>12</sup></td><td>300</td></tr><tr><td>H<sup>13</sup>CO</td><td>1.5 × 10<sup>15</sup></td><td>1.5 × 10<sup>14</sup></td><td>700</td></tr><tr><td>H<sub>2</sub>O</td><td>5.0 × 10<sup>11</sup></td><td>1.5 × 10<sup>12</sup></td><td>300</td></tr><tr><td>H<sub>2</sub>CCO</td><td>1.5 × 10<sup>15</sup></td><td>2.0 × 10<sup>14</sup></td><td>700</td></tr><tr><td>H<sub>2</sub>O, v<sub>2</sub></td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td>H<sup>13</sup>CN, v<sub>2</sub> = 1</td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td>HC<sub>3</sub>N</td><td>2.0 × 10<sup>17</sup></td><td>5.0 × 10<sup>14</sup></td><td>700</td></tr><tr><td>HC<sub>3</sub>N, v = 0</td><td>4.0 × 10<sup>17</sup></td><td>6.0 × 10<sup>14</sup></td><td>700</td></tr><tr><td>HCN, v<sub>2</sub> = 1</td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td>HCN, v<sub>2</sub> = 2</td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td>HD<sup>18</sup>O</td><td>5.0 × 10<sup>11</sup></td><td>3.0 × 10<sup>12</sup></td><td>300</td></tr><tr><td>HDCO</td><td>7.0 × 10<sup>12</sup></td><td>9.0 × 10<sup>12</sup></td><td>300</td></tr><tr><td>HN<sup>13</sup>C</td><td>8.0 × 10<sup>11</sup></td><td>1.0 × 10<sup>12</sup></td><td>300</td></tr><tr><td>HN<sup>13</sup>CO</td><td>5.0 × 10<sup>12</sup></td><td>1.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td>HNC, v<sub>2</sub> = 1</td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td>HNCO</td><td>5.0 × 10<sup>12</sup></td><td>1.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td>NH<sub>3</sub>CHO</td><td>3.0 × 10<sup>13</sup></td><td>3.0 × 10<sup>13</sup></td><td>700</td></tr><tr><td>NH<sub>2</sub>D</td><td>4.0 × 10<sup>12</sup></td><td>8.0 × 10<sup>12</sup></td><td>300</td></tr><tr><td>NH<sub>3</sub>, v<sub>2</sub></td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td>NS</td><td>1.5 × 10<sup>13</sup></td><td>1.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td>O<sub>2</sub></td><td>2.0 × 10<sup>17</sup></td><td>4.0 × 10<sup>17</sup></td><td>700</td></tr><tr><td>OCS</td><td>3.0 × 10<sup>17</sup></td><td>3.0 × 10<sup>15</sup></td><td>700</td></tr><tr><td>OD</td><td>1.0 × 10<sup>14</sup></td><td>1.5 × 10<sup>13</sup></td><td>700</td></tr><tr><td>OH</td><td>3.0 × 10<sup>15</sup></td><td>1.5 × 10<sup>14</sup></td><td>300</td></tr><tr><td>SiS</td><td>9.0 × 10<sup>14</sup></td><td>6.0 × 10<sup>13</sup></td><td>700</td></tr><tr><td>SO<sub>2</sub>, v<sub>2</sub> = 2</td><td>N/A</td><td>N/A</td><td>700</td></tr><tr><td><sup>33</sup>SO</td><td>3.0 × 10<sup>13</sup></td><td>1.5 × 10<sup>13</sup></td><td>700</td></tr><tr><td><sup>34</sup>SO</td><td>5.0 × 10<sup>13</sup></td><td>2.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td><sup>33</sup>SO<sub>2</sub></td><td>2.0 × 10<sup>13</sup></td><td>4.0 × 10<sup>13</sup></td><td>300</td></tr><tr><td><sup>34</sup>SO<sub>2</sub></td><td>1.5 × 10<sup>13</sup></td><td>3.0 × 10<sup>13</sup></td><td>200</td></tr></tbody></table>

<table><caption>**Table 9**

Column Density Upper Limit for Unmodeled Species in Orion-S</caption><thead><tr><th>Species</th><th>N<sub>t</sub>, T<sub>ex</sub>=40 K</th><th>N<sub>t</sub>, T<sub>ex</sub>=80 K</th></tr></thead><tbody><tr><td>CH<sub>3</sub>OCH<sub>3</sub></td><td>1.5 × 10<sup>14</sup></td><td>3 × 10<sup>14</sup></td></tr><tr><td>CO<sup>+</sup></td><td>2 × 10<sup>12</sup></td><td>1.5 × 10<sup>12</sup></td></tr><tr><td>H<sub>2</sub><sup>18</sup>O</td><td>2 × 10<sup>12</sup></td><td>3 × 10<sup>12</sup></td></tr><tr><td>HC<sup>15</sup>N</td><td>1 × 10<sup>12</sup></td><td>2 × 10<sup>12</sup></td></tr><tr><td>HCS<sup>+</sup></td><td>2.5 × 10<sup>13</sup></td><td>1 × 10<sup>13</sup></td></tr><tr><td>NH<sub>2</sub></td><td>7 × 10<sup>12</sup></td><td>1 × 10<sup>13</sup></td></tr><tr><td>SH<sup>+</sup></td><td>4 × 10<sup>12</sup></td><td>7 × 10<sup>12</sup></td></tr><tr><td>SiO</td><td>1 × 10<sup>13</sup></td><td>4 × 10<sup>12</sup></td></tr></tbody></table>
SO and $SO_2$ are often associated with shocked gas, they do not appear in Figure 18. This is because these molecules were optically thick and Crockett et al. (2014) did not provide column densities.

The best agreement with molecular abundances in the narrow component of Orion-S is with the Extended Ridge of Orion-KL (Figure 19), where we obtain $\langle \frac{X_{KL}}{X_S} \rangle = 7$ (SD = 14).
For the broad component (dashed line) we obtain $\langle \frac{X_{KL}}{X_S} \rangle = 0.3$ (SD = 0.5).

Given that the best match to the abundances in the narrow component of Orion-S is the ER of Orion-KL, it seems as though these species/components do indeed trace quiescent gas. In particular, it probably is the same gas out of which both star-forming regions have been formed. The broad component of Orion-S, however, seems better matched to the CR and Plateau of Orion-KL. Figures 16 to 19 only have a few broad component points and, therefore, it is difficult to make any strong statistical arguments based on these data alone. However, this evidence along with the chemical abundance analysis presented in Section 3.3.1 provide fairly strong support for the idea that shocks have also had an influence on the chemistry of Orion-S.

### 3.3.3. Species Detected in Orion-KL but Not in Orion-S

It is well known that Orion-KL has an incredibly rich molecular chemistry (e.g., Schilke et al. 1997a; Comito et al. 2005; Tercero et al. 2005; Leurini et al. 2006; Olofsson et al. 2007; Persson et al. 2007). However, it is possible that the species that were detected in Orion-KL, but not in Orion-S, exist in the latter source, but at levels too weak to be detected. Some of these species might be observable with ALMA at lower frequencies. In this section, we explore this possibility by providing upper limits for the abundances of all the species detected in Orion-KL by Crockett et al. (2014) but not detected in Orion-S.

Modeling was accomplished by fixing $T_{ex}$, $\Delta V_{FWHM}$, $\Omega$, and $V_{LSR}$ and finding the column density that produced transitions for which the intensities were <3$\sigma$ across all the HEXOS bands. Since we were able to model all of our species using, at most, two components, we determine two different column density upper limits: one assuming that the undetected emission arises from the narrow component and assuming that it originates in the broad component. For the narrow component, we used fixed values of $T_{ex} = 40$ K, $\Delta V_{FWHM} = 4$ km s$^{-1}$, $\Omega = 60''$, and $V_{LSR} = 7.1$, which were found to be typical for the narrow component (see Tables 5 and 6). For the broad component, we used fixed values of $T_{ex} = 80$ K, $\Delta V_{FWHM} = 8$ km s$^{-1}$, $\Omega = 40''$, and $V_{LSR} = 7.1$, which were found to be typical for the broad component (see Table 6). Results are listed in Table 8. Column 1 is the species name, column 2 is the upper limit total column density assuming the gas arises in the narrow component, column 3 is the upper limit total column density assuming the gas originates in the broad component, and column 4 lists the maximum upper state energy for the model (i.e., no transitions with $E > E_{up}$ were modeled). Different values of $E_{up}$ were used for different molecules to keep the number of modeled transitions to a reasonable value.

There are also a number of species that we detected ($S/N > 3\sigma$) but did not model since their $S/Ns$ were <5$\sigma$. Using the same assumptions as for the undetected species, we provide upper limits for the abundances for these species in

Table 9. The column density limits listed in Table 9 are approximately an order of magnitude or more smaller than the column densities of the same species detected in Orion-KL (Crockett et al. 2014).

A possible question is whether the species listed in Table 8 could actually exist in Orion-S even though they are undetected. We inspect the excitation conditions of these species, examining $E_{up}$ of all possible transitions, to determine whether they are detectable given the reported noise of the HIFI bands. A portion of these species have $E_{up}$ much greater than 500 K (e.g., H$_2$O $v_2$, H$^{13}$CN $v_2 = 1$, HC$_3$N, HC$_3$N $v = 0$, HCN $v_2 = 1$, HCN $v_2 = 2$, OCS, etc.). Based on the analysis done in this paper, Orion-S can barely excite species with $E_{up} > 500$ K. Therefore, transitions of these species would not be observed above the noise, even if they were present. In addition, there are some complex organic species with transition of $E_{up} < 100$ K which we also did not detect (e.g., CH$_2$NH, CH$_2$DOH, CH$_2$NH, CH$_3$OCHO, etc.). This is not surprising given the absence of hot-core chemistry in Orion-S.

## 4. SUMMARY AND CONCLUSION

We have presented results from a comprehensive spectral survey toward Orion South, taken with the HIFI instrument on board the Herschel Space Observatory covering the frequency range of 480 to 1900 GHz with a resolution of 1.1 MHz. We detected 685 spectral lines with $S/Ns > 3\sigma$ originating from 52 different molecular and atomic species. Using the CASSIS spectral line analysis software package, we modeled each of the detected species assuming conditions of Local Thermodynamic Equilibrium. Based on this modeling, we found evidence for three different cloud components: a cool ($T_{ex} \sim 20-40$ K), spatially extended ($>60''$), and quiescent ($\Delta V_{FWHM} \sim 4$ km s$^{-1}$) component; a warmer ($T_{ex} \sim 80-100$ K), less spatially extended ($\sim 30''$), and more dynamic ($\Delta V_{FWHM} \sim 8$ km s$^{-1}$) component, which is likely affected by embedded outflows; and a kinematically distinct region dominated by emission from species that trace UV irradiation. Indirect evidence to support the existence of the first two components can be inferred from McMullin et al. (1993) who mapped the region in a few spectral lines (SiO, H$^{13}$CO$^+$, SO$_2$, CH$_3$OH, and HC$_3$N) with the BIMA array. Their H$^{13}$CO$^+$ and HC$_3$N data confirm the existence of a fairly extended ($\sim 1'$) quiescent (FWHM $\sim 3$ km s$^{-1}$) component, whereas the SO$_2$ and CH$_3$OH data reveal a smaller emitting region ($\sim 20''$) of warm gas ($\sim 75$ K). While the spectra for the latter two species are too weak to determine line widths, their SiO data reveal a similarly small region (offset by only a few arc seconds from the SO$_2$ and CH$_3$OH emission peaks) with broad-line widths ($\sim 7$ km s$^{-1}$). In addition, McMullin et al. (1993) reports column densities of SO$_2$ and H$^{13}$CO$^+$ of $<2 \times 10^{-10}$ and $4 \times 10^{-11}$, respectively, which compare favorably to the values reported in Table 7. Finally, while there are no higher resolution observations to confirm the existence of the third component (i.e., the UV irradiated region), since CO$^+$ is only ever detected in PDRs, its presence in our data strongly suggests that such a component must exist.

We also presented a comprehensive chemical abundance comparison between Orion-KL and Orion-S; two star-forming regions that potentially formed from the same natal molecular gas but are at different evolutionary stages. Based on a paucity of complex molecules in Orion-S, we found little chemical evidence for the existence of a significant “hot-core”
component. This is likely due to the fact that the hot cores associated with the embedded star formation have either not had sufficient time to develop chemically or are simply too small for their line emission to be detected in the large Herschel beam, or Orion-S is not a massive star-forming region and hot cores massive enough to produce their characteristic rich spectra simply do not exist. The presence of a number of UV tracers such as [C II], [C I], CH, CH+, SH+, and CO+, and the fact that transitions of these species have velocities that are $1-1.5 \text{ km s}^{-1}$, higher than those of the quiescent gas, suggest that these species arise from a kinematically distinct PDR; most likely the UV illuminated surface of the cloud. The best match to the chemical abundances in the cooler, quiescent gas in Orion-S is with the quiescent extended ridge of Orion-KL, indicating that most of the gas in Orion-S is still quiescent as well, and relatively unaffected by higher temperature or UV driven chemistry. The best agreement with the warmer, broad component of Orion-S is with the Orion-KL Plateau and Compact Ridge regions, suggesting that shocks have had an influence on the overall chemistry in Orion-S.

HIFI was designed and built by a consortium of institutes and university departments from across Europe, Canada, and the United States under the leadership of the SRON Netherlands Institute for Space Research, Groningen, The Netherlands and with major contributions from Germany, France, and the US. Consortium members are Canada: CSA, U. Waterloo; France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland, NUI Maynooth; Italy: ASI, IFSI-INAF, Osservatorio Astrofisico di Arcetri- INAF; Netherlands: SRON, TUD; Poland: CAMK, CBK; Spain: Observatorio Astronmico Nacional (IGN), Centro de Astrobiologa (CSIC-INTA). Sweden: Chalmers University of Technology—MC2, RSS & GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm University—Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA: Caltech, JPL, NHSC. We also need to acknowledge the support by the Deutsche Forschungsgemeinschaft (DFG) via the collaborative research grant SFB 956, project C1 & C3, as well as the ERC and the Spanish MINECO for funding support under grants ERC-2013-Syg-610256 and AYA2012-32032. Support for this work was provided, in part, by a National Sciences and Engineering Research Council of Canada (NSERC) grant to R. Plume and K. Tahani and by NASA through an award issued by JPL/Caltech. This work was carried out in part at the Jet Propulsion Laboratory, which is operated for NASA by the California Institute of Technology.

Facility: Herschel.

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