Chemical Complexity of Phosphorous-bearing Species in Various Regions of the Interstellar Medium

Phosphorus (P) and its compounds play a crucial role in the chemical evolution in galaxies. Phosphorus is an essential element of life and is one of the main biogenic elements. Its origin in the terrestrial system is yet to be fully understood. The Atacama Large Millimeter/submillimeter Array (ALMA) and European Space Agency Probe Rosetta suggest that P-related species might have traveled from star-forming regions to Earth (e.g., Altwegg et al. 2016; Rivilla et al. 2020) through comets. The interstellar chemistry of P-bearing molecules has significant astrophysical relevance, and very little has been revealed so far. P-bearing compounds are the major components of any living system, where they carry out numerous biochemical functions. P-bearing molecules play a significant role in forming large biomolecules or living organisms; they store and transmit the genetic information in nucleic acids and work in nucleotides as precursors in the synthesis of RNA and DNA (Maciá et al. 1997). Moreover, these molecules are the essential components of phospholipids (main characteristic features of cellular membranes; Maciá 2005).

Phosphorus is relatively rare in the interstellar medium (ISM). However, it is ubiquitous in various meteorites (Jarosewich 1990; Lodders 2003; Pasek 2019). On average, P is the 13th most abundant element in a typical meteoritic material and the 11th most abundant element in Earth's crust (Maciá 2005).

Jura & York (1978) identified P⁺ with a cosmic abundance of ∼2 × 10⁻⁷ in hot regions (∼1200 K). P-bearing molecules such as PN, PO, HCP, CP, CCP, and PH₃ have been observed in circumstellar envelopes around evolved stars (Guelin et al. 1990; Agúndez et al. 2007, 2008, 2014; Tenenbaum et al. 2007; Halfen et al. 2008; Milam et al. 2008; Tenenbaum & Ziurys 2008; De Beck et al. 2013; Ziurys et al. 2018). PN was detected in several star-forming regions (Turner & Bally 1987; Ziurys 1987; Turner et al. 1990; Caux et al. 2011; Yamaguchi et al. 2011; Fontani et al. 2016, 2019; Mininni et al. 2018), and it remained the only P-containing species identified in the dense ISM for many years (Ziurys 1987; Turner et al. 1990; Fontani et al. 2016). Rivilla et al. (2016) reported the first detection of PO, along with PN, toward two massive star-forming regions, W51 1e/2e and W3(OH), by using the IRAM 30 m telescope. Thereafter, both PO and PN were simultaneously observed in several low- and high-mass star-forming regions and the Galactic center (Lefloch et al. 2016; Rivilla et al. 2018, 2020; Bergner et al. 2019; Bernal et al. 2021).

Theoretical investigation on P-chemistry modeling has been extensively reported in several studies (Millar 1991; Charnley & Millar 1994; Aota & Aikawa 2012; Lefloch et al. 2016; Rivilla et al. 2016; Jiménez-Serra et al. 2018). However, the P-chemistry of the dense cloud region is not well constrained. The main uncertainty lies in the depletion factor of the initial elemental abundance of P. The degree of complexity of gas-phase abundance within the gas and the exact depletion of different elements onto the grains are still uncertain (Jenkins 2009). Very recently, Nguyen et al. (2020) discovered the chemical desorption of phosphine. This kind of study is indeed very crucial in constraining the modeling parameters.

Phosphine (PH₃) is a relatively stable molecule that could hold a significant fraction of P in various astronomical environments. Scientists consider PH₃ to be a biosignature (Sousa-Silva et al. 2020). PH₃ was identified in the planets of our solar system containing a reducing atmosphere. The Voyager data confirmed PH₃ within Jupiter and Saturn with volume mixing ratios of 0.6 and 2 ppm, respectively (Maciá 2005). It further suggested that the P/H ratio shows harmony with the solar value. Fletcher et al. (2009) derived the global distribution of PH₃ on Jupiter and Saturn using 2.5 cm⁻¹ spectral resolution of Cassini/CIRS observations. PH₃ could be produced deep inside the reducing atmospheres of giant planets (Bregman et al. 1975; Tarrago et al. 1992) at high temperatures and pressures and dredged upward by the convection (Noll & Marley 1997; Visscher et al. 2006). Very surprisingly, in Venus, there is no such reducing atmosphere, but ∼20 ppb of PH₃ was inferred (Greaves et al. 2021) in the deck of Venus's atmosphere. They could not account for such a high amount of PH₃ with the steady-state chemical models, including the photochemical pathways. They also explored the other abiotic routes for explaining this high abundance. But none of them seem to be suitable. They speculated some unknown photo- or geochemical origin of it. More far-reaching, its high abundance by some of the biological means could not be ruled out. Bains et al. (2020) also could not explain the presence of PH₃ in Venus's clouds by any abiotic mechanism based on their state-of-the-art understanding. This discovery opened up a series of debates. Very recently, Villanueva et al. (2020) questioned the analysis and interpretation of the spectroscopic data used in Greaves et al. (2021). These authors claimed that there is no PH₃ in Venus's atmosphere. Snellen et al. (2020) also found no statistical evidence for PH₃ in the atmosphere of Venus.

Among the simple P-bearing species, PH₃ was tentatively identified (J = 1 → 0, 266.9 GHz) in the envelope of the carbon-rich stars IRC +10216 and CRL 2688 (Agúndez et al. 2008; Tenenbaum & Ziurys 2008). Later, the presence of PH₃ was confirmed by Agúndez et al. (2014). They observed the J = 2 → 1 rotational transition of PH₃ (at 534 GHz) in IRC +10216 using the HIFI instrument on board Herschel. They predicted a very high abundance of PH₃ in this region. Suppose that the gas-phase reactions are responsible for such a high abundance of PH₃. In that case, it should occur through a barrierless reaction or with a reaction with a little barrier at low temperatures. Here, we employ various state-of-the-art chemical models to understand the formation of PH₃ under different interstellar conditions like diffuse cloud, interstellar photon-dominated regions or photodissociation regions (PDRs), and hot core regions.

Recently, Chantzos et al. (2020) attempted to observe HCP (21), CP (21), PN (21), and PO (21) with the IRAM 30 m telescope along the line of sight to the compact extragalactic quasar B0355+508. They were unable to detect these transitions along this line of sight. But based on their observations, they predicted 3σ upper limits of these transitions. However, they successfully identified HNC (10), CN (10), and C³⁴S (21) in absorption and ¹³CO (10) in emission along the same line of sight.

We organize this paper as follows. In Section 2, we prepare a currently available chemical network for the P-bearing species, and in Section 3, we present their various kinetic information. Section 4 is devoted to the chemical model explaining P-chemistry in different parts of the ISM. In Section 5, we describe our results obtained with the radiative transfer model for the diffuse and hot core region. In Section 6, we explain the infrared spectrum of PH₃ ice under various circumstances, and finally, in Section 7, we draw our conclusion.

We prepare an extensive network to study the chemistry of the P-bearing species. Since nitrogen (N) and P have five electrons in the valance shell, P-chemistry is sometimes considered analogous to N-chemistry. For example, successive hydrogenation of N yields NH₃, whereas for P it yields PH₃. However, the presence of NH₃ is ubiquitous in the ISM (Cheung et al. 1968; Wilson & Pauls 1979; Keene et al. 1983; Mauersberger et al. 1988), but the presence of PH₃ is not so universal (Thorne et al. 1983). On the other hand, P's chemistry can notably differ from the N-related chemistry under the physical conditions prevailing in the different star-forming regions. In this paper, we prepare an extensive network of P by following the chemical pathways explained in Thorne et al. (1984), Adams et al. (1990), Millar (1991), Fontani et al. (2016), Jiménez-Serra et al. (2018), Sousa-Silva et al. (2020), Rivilla et al. (2016), and Chantzos et al. (2020). Charnley & Millar (1994) and Anicich (1993) differentiated the formation/destruction mechanism of PH_n (n = 1, 2, 3) and their cationic species. Thorne et al. (1984) presented the reaction pathways for P, PO, P⁺, PO⁺, PH⁺, HPO⁺, and H₂PO⁺. Furthermore, in a recent study, Chantzos et al. (2020) extend the gas-phase chemical network of PN and PO in accordance with Millar et al. (1987) and Jiménez-Serra et al. (2018). We use the kinetic data of chemical reactions from the KInetic Database for Astrochemistry (KIDA; Wakelam et al. 2015) and the UMIST Database for Astrochemistry (UDfA; McElroy et al. 2013). In Table A1, we show all the reactions considered in our P-network.

A continuous exchange of chemical ingredients within the gas and grain can determine the chemical complexity of the ISM. The major drawback in constraining this chemical process by astrochemical modeling is the lack of information about the interaction energy of the chemical species with the grain surface. The fate of the chemical complexity on the grain surface depends on the residence time of the incoming species on the grain. Thus, the binding energy (BE) of interstellar species plays a crucial role in the chemical complexity of the ISM.

In a periodic treatment of surface adsorption phenomena, BE is related to the interaction energy (ΔE) as

$$\begin{eqnarray}&&\mathrm{BE}=-{\rm{\Delta }}E.\end{eqnarray} \tag{ 1 }$$

For a bounded adsorbate BE is a positive quantity and is defined as

$$\begin{eqnarray}&&\mathrm{BE}=({E}_{\mathrm{surface}}+{E}_{\mathrm{species}})-{E}_{\mathrm{ss}},\end{eqnarray} \tag{ 2 }$$

where E_ss is the optimized energy for the complex system where a species is placed at a distance from the grain surface through a weak van der Waals interaction. E_surface and E_species are the optimized energies of the grain surface and species, respectively.

In dense molecular clouds, a sizable portion of the grain mantle would cover water, methanol, and CO molecules. In denser regions, gas-phase H₂ could easily accrete on the grain and mostly gets back to the gas phase owing to its low sticking probability and lower BEs. But some H₂ could be trapped under some accreting species. For example, one H₂ may accrete on another H₂ before it is desorbed. In this situation, the "encounter desorption" mechanism is supposed to be an essential means of transportation of H₂ to the gas phase (Hincelin et al. 2015; Das et al. 2021). Usually, in the literature, only the encounter desorption of H₂ is considered owing to its wide presence in the dense ISM. Recently, Chang et al. (2021) considered the encounter desorption of H atom as well. But the BEs of the interstellar species with the different substrates are unknown. Looking at these aspects, in Table 1 we report the BEs of some relevant P-bearing species by considering H₂O, CO, CH₃OH, and H₂ as a substrate. The computed BEs should have a range of values that depend on the molecule's position on the ice (Sil et al. 2017; Das et al. 2018; Ferrero et al. 2020). Whenever we have obtained different BE values at other binding sites, we exert the average of some of our calculations and note them in Table 1. In this paper, we only consider the BE with water substrate for our simulation. BEs with the other substrates are provided for future usage. We perform all the BE calculations by the Gaussian 09 suite of programs (Frisch et al. 2013) using Equations (1) and (2). To find the optimized energy of all structures, we use a second-order Møller–Plesset (MP2) method with an aug-cc-pVDZ basis set (Dunning 1989). We did not include the zero-point vibrational energy (ZPVE) and basis set superposition error (BSSE) correction (Das et al. 2018). A fully optimized ground-state structure is verified to be a stationary point (having nonimaginary frequency) by harmonic vibrational frequency analysis. Ground-state spin multiplicity of the species is also noted in Table 1. To evaluate it, we take the help of the Gaussian 09 suite of programs and run separate calculations (job type "opt + freq"), each with different spin multiplicities, and then compare the results between them. The lowest-energy electronic state solution of the chosen spin multiplicity is the ground state noted.

Due to the similarities between the structures of NH₃ and PH₃, Chantzos et al. (2020) considered the BE of PH₃ to be the same as that of NH₃. The BE of NH₃ is 5800 K according to the KIDA database. Since they also used this BE of PH₃, they did not obtain much PH₃ by the thermal desorption process in the cold regions. They pointed out that most of the PH₃ came to the gas phase by photodesorption rather than thermal desorption from the colder region. In our calculation, we found a noticeable decrease in the BE of PH₃, which could enable PH₃ to populate the gas phase by thermal desorption even at low temperatures. Das et al. (2018) noted BE of NH₃ with a c-hexamer configuration of water ∼5163 K, whereas we have found it to be ∼2395 K for PH₃. With significant ice constituents (monomer of CO, CH₃OH, H₂O, and H₂), our computed BE of PH₃ seems to be <1000 K. Sousa-Silva et al. (2020) noted that PH₃ has very low water solubility (at 17°C only 22.8 ml of gaseous PH₃ dissolves in 100 ml of water), and it does not easily stick to aerosols. The recent claim of PH₃ in the Venusian atmosphere (Greaves et al. 2021) prompted us to determine the BE of PH₃ with some other species such as sulfuric acid (H₂SO₄) and benzene (C₆H₆). These are the principal constituents of the Venus atmosphere. Our computed BE with H₂SO₄ and benzene is ∼2271 and ∼3094 K, respectively. However, Snellen et al. (2020) and Villanueva et al. (2020) questioned the identification of Greaves et al. (2021). Table 1 also provides our calculated enthalpy of formation for several P-bearing species and the BE values noted in various places.

We use the reaction pathways shown in Appendix A (see Table A1) to check the fate of the P-bearing species in various parts of the ISM (diffuse clouds, PDRs, hot corino, and hot cores). Here, we employ mainly two models to study the chemical evolution of these species: (a) the spectral synthesis code Cloudy (version 17.02, last described by Ferland et al. 2017) and (b) the Chemical Model for Molecular Cloud (hereafter CMMC) code (Das et al. 2015; Gorai et al. 2017a, 2017b, 2020; Sil et al. 2018; Shimonishi et al. 2020; Das et al. 2021).

4.1. Spectral Synthesis Code

We use a photoionization code, Cloudy, which simulates matter under a broad range of interstellar conditions. It is for general use under an open-source license. ¹¹ Using the Cloudy code, we check the fate of P-bearing species in the diffuse cloud (Das et al. 2020) and PDR environment.

4.1.1. Diffuse Cloud Model

Based on the observations carried out by Chantzos et al. (2020), they prepared a diffuse cloud model to explain the observed abundance of HNC, CN, CS, and CO. We also employ a similar process to explain the observed abundances of the four molecules and then look at the fate of the P-bearing molecules under these circumstances. For this modeling, we consider the initial elemental abundance as given in Chantzos et al. (2020). But in Cloudy, only the atomic elemental abundance is allowed (no ionic or molecular form), so we use these abundances as the initial elemental abundance (see Table 2). For each calculation, we check whether the microphysics considered is time steady or not. We notice that the largest reaction timescale is much shorter than the diffuse cloud's typical lifetime, ∼10⁷ yr. We use the grain size distribution, which is appropriate for the R_V = 3.1 extinction curve of Mathis et al. (1977). This grain size distribution is called "ISM" in Cloudy. We use the anisotropic radiation field, which is appropriate for the ISM to specify the intensity of the incident local interstellar radiation field. Additionally, we include the cosmic-ray microwave background with a redshift value of 1.52 (Wenger et al. 2000).

Table 2. Initial Elemental Abundance for the Diffuse Cloud and PDR Model Considered in the Cloudy Code

Element	Abundance	Element	Abundance
H	1.0	Si	3.2 × 10−5
He	8.5 × 10−2	Fe	3.2 × 10−5
N	6.8 × 10−5	Na	1.7 × 10−6
O	4.9 × 10−4	Mg	3.9 × 10−5
C	2.7 × 10−4	Cl	3.2 × 10−7
S	1.3 × 10−5	P	2.6 × 10−7
		F	3.6 × 10−8

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To check the validity of our model against the observed abundances, we first explain the observed abundance of CO, CN, CS, and HNC as obtained in Chantzos et al. (2020). For this, we use a parameter space for the formation of these species. Our parameter space consists of a density variation of 100–600 cm⁻³ and a cosmic-ray ionization rate of H₂ (${\zeta }_{{{\rm{H}}}_{2}}$) in between 10⁻¹⁷ and 10⁻¹⁴ s⁻¹. We use a stopping criterion at different A_V values. The cases obtained by using stopping criterion A_V = 1–5 mag are given in Figure 1. The observed abundances of CO, CN, CS, and HNC of Chantzos et al. (2020) toward the cloud having a V_LSR = −17 km s⁻¹ are highlighted by the contours. Figure 2 shows the variation of temperature at the end of the calculation for A_V = 1–3 mag. The parameter space also consists of the variation of n_H and ${\zeta }_{{{\rm{H}}}_{2}}$. Temperature variation is shown with the color bar.

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From Figures 1 and 2, we see that the parameters that were adopted by Chantzos et al. (2020) for the diffuse cloud model (i.e., ${\zeta }_{{{\rm{H}}}_{2}}=1.7\times {10}^{-16}\,{{\rm{s}}}^{-1}$, n_H = 300 cm⁻³, and T_gas = 40 K) can also reproduce the observed abundances of CO, CN, CS, and HNC with our model. Figure 3 shows the obtained abundances of the four molecules when we consider n_H = 300 cm⁻³ and ${\zeta }_{{{\rm{H}}}_{2}}$ = 1.7 × 10⁻¹⁶ s⁻¹. The X-axis shows the visual extinction of the cloud, and the Y-axis shows the abundance of n_H. Our results are in agreement with the observation between A_V = 2 and 3 mag. We highlight the best-suited zone by the black dashed curve. Thus, based on Figures 1–3, we use A_V = 2 mag, ${\zeta }_{{{\rm{H}}}_{2}}=1.7\,\times \,{10}^{-16}$, and n_H = 300 cm⁻³ as the best-fitted parameters to explain the observed abundances of these species. This yields A_V /N(H) = 5.398 × 10⁻²² mag cm² and a dust-to-gas ratio of 6.594 × 10⁻³. Table 3 compares our obtained optical depth and column densities with the observations. Obtained optical depths of CN and HNC are close to the observed values, but the column densities diverge by a few factors. This slight mismatch is because the Cloudy model deals with steady-state values, but the column densities are time dependent in reality.

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Table 3. Estimated Column Density and Optical Depth of the Observed Molecules for the Diffuse Cloud Model Obtained with the Cloudy Code

Species	Transitions	Eup (K)	Frequency (GHz)	Optical Depth (τ)		Total Column Density (cm−2)
				Model	Observation a	Model	Observation a
CN	N = 1–0, J = 1/2–1/2, F = 3/2–1/2	5.5	113.16867	0.256522	0.23 ± 0.07	1.47 × 1012	(0.87 ± 0.28) × 1013
HNC	J = 1–0	4.4	90.66357	0.33962	0.50 ± 0.10	2.11 × 1011	(0.69 ± 0.16) × 1012
C34S	J = 2–1	6.9	96.41295	⋯	0.04 ± 0.02	⋯	(1.64 ± 0.82) × 1011
13CO	J = 1–0	5.3	110.20135	⋯	0.154 ± 0.004	⋯	(3.98 ± 0.16) × 1014
HCP	J = 2–1	5.8	79.90329	2.56 × 10−6	<0.02	1.10 × 108	<2.27 × 1012
PN	J = 2–1	6.8	93.97977	0.238015	<0.02	1.62 × 1011	<4.20 × 1010
CP	N = 2–1, J = 3/2–1/2, F = 2–1	6.8	95.16416	4.92 × 10−4	<0.02	1.16 × 1010	<1.26 × 1012
PO	J = 5/2–3/2, Ω = 1/2, F = 3–2, e	8.4	108.99845	0.0138432	<0.02	9.08 × 1010	<4.29 × 1011
PO	J = 5/2–3/2, Ω = 1/2, F = 2–1, e	8.4	109.04540	0.0125113	<0.02	9.08 × 1010	<6.70 × 1011
PO	J = 5/2–3/2, Ω = 1/2, F = 3–2, f	8.4	109.20620	0.0055508	<0.02	9.08 × 1010	<4.34 × 1011
PO	J = 5/2–3/2, Ω = 1/2, F = 2–1, f	8.4	109.28119	0.0182699	<0.02	9.08 × 1010	<6.69 × 1011

Note.

^a Chantzos et al. (2020).

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Figure 4 depicts the abundances of most of the P-bearing species considered in this study. The left panel shows the neutral P-bearing species, whereas the right panel shows the P-bearing ions. It is interesting to note that although we consider a neutral atomic abundance in our model, the abundance of P⁺ is high owing to the strong cosmic-ray ionization and presence of a nonextinguished interstellar radiation field. The abundances of the comparatively larger P-bearing species are very low in the diffuse environment, and it would be pretty challenging to observe them. We notice that some simple neutral P-bearing species like PN, PO, HCP, CP, and PH are significantly more abundant than PH₃. The abundances of PO and PN appear to be comparable to each other (with PO/PN < 1 when A_V ≥ 2). Table 3 shows that the obtained column densities for this case are listed. There is an excellent agreement between the observed and the modeled optical depths of CN and HNC with the Cloudy code.

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In Figure 5, the abundances of H, H₂, and N are shown along with the major hydrogen-related ions, H⁺, ${{{\rm{H}}}_{2}}^{+}$, and ${{{\rm{H}}}_{3}}^{+}$. In Cloudy, the cloud temperature is calculated self-consistently, depending on the known excitations. The primary source of such excitation is cosmic-ray ionization. Figure 5 depicts the gas temperature variation with A_V with the solid red curve. Deep inside the cloud, the temperature drops sufficiently and enhances the formation of molecular hydrogen. For A_V = 2 mag, we obtained a temperature ∼10 K when we use a cosmic-ray ionization rate of 1.7 × 10⁻¹⁶ s⁻¹. It is a little lower than that used in Chantzos et al. (2020).

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4.1.2. PDR Model

Here, we use the Cloudy code to explain the abundances of the P-bearing species in the PDRs. The Cloudy code is capable of considering realistic physical conditions around this region. PDRs play an essential role in interstellar chemistry. They are responsible for the emission characteristics of ISM and dominate the infrared and submillimeter spectra of star-forming regions and galaxies. In the PDRs, H₂-nonionizing far-ultraviolet (FUV) photons from stellar sources control gas heating and chemistry. Here, the gas is heated by the FUV radiation (6 eV < hν < 13.6 eV, from the ambient UV field and hot stars) and cooled via the emission of spectral line radiation of atomic and molecular species and continuum emission by dust. Cloudy includes various PDR benchmark models (Röllig et al. 2007), giving a range of physical conditions associated with PDR environments. We chose the F4 model of Röllig et al. (2007). This model considered a constant gas temperature of 50 K and dust temperature of 20 K. A plane-parallel, semi-infinite cloud with total constant density of 10^5.5 cm⁻³ and standard UV field as χ = 10⁵ times the Draine (1978) field (χ = 1.71 G₀) is used. The cosmic-ray H ionization rate ζ = 5 × 10⁻¹⁷ s⁻¹, the visual extinction A_V = 6.289 ×10⁻²² N_H,tot, and a dust-to-gas ratio of 7.646 × 10⁻³ are assumed. We also consider these physical parameters along with the initial elemental abundances as considered in our diffuse cloud model (Table 2). Figure 6 shows the abundances of important P-bearing species. In the diffuse cloud region (see Figure 4), we obtained PO and PN comparable abundances. However, for the PDR model, we found a very high abundance of PN compared to PO. Its peak abundance appears to be higher by several orders of magnitude than PO (i.e., PO/PN ≪ 1). We use the photodestruction rate of PO and PN from Jiménez-Serra et al. (2018). The photodissociation rate of PO is estimated from the photodissociation rate of NO. Similarly, the photodissociation rate of PN is estimated from the photodissociation rate of N₂. These values reflect a higher photodissociation rate of PO (with α = 3 × 10⁻¹⁰ s⁻¹ and γ = 2.0) than PN (with α = 5 × 10⁻¹² s⁻¹ and γ = 3.0). As in Jiménez-Serra et al. (2018), here we also find that even if PN is photodissociated, it may be able to revert again via P + CN → PN + C owing to the presence of a large amount of atomic P in the gas phase. The rapid conversion of PO to PN by the reaction PO + N → PN + O is also crucial for maintaining a higher abundance of PN than PO.

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It is not straightforward to relate the diffuse cloud region cases with the PDR because of different physical circumstances. In the diffuse cloud model, we have obtained the peak abundances of PN and PO at A_V = 2–3 mag, whereas in the PDR model we have received these peaks around A_V = 6–7 mag. In both cases, the PO/PN ratio is <1, but this ratio seems to be ≪1 for the PDR. Jiménez-Serra et al. (2018) also modeled the effect of intense UV photon illumination by varying the interstellar radiation field within ranges typical of PDRs. They showed for higher extinctions (A_V = 7.5 mag) and under high-UV radiation fields (χ = 10⁴ Habing) that the abundance of PN always remains above that of PO for both long-lived and short-lived collapse stages. Figure 7 shows the abundances of H, H⁺, H₂, ${{{\rm{H}}}_{2}}^{+}$, ${{{\rm{H}}}_{3}}^{+}$, N, and electrons.

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4.2. CMMC Code

In the earlier section, we have implemented a spectral synthesis code, Cloudy, to study the chemical evolution of the P-bearing species in the diffuse cloud region. Comparatively larger P-bearing species are not very profuse in space, which creates a burden constraining the understanding of the P-bearing species of the ISM. The major drawback in the P-chemistry modeling is the uncertainty of the P depletion factor. The grain-surface chemistry plays a significant role in shaping the chemical complexity in these regions. Since Cloudy only considers the surface reactions of some key species, it would be impractical to apply the Cloudy code in the dense cloud region. Thus, we use our gas-grain CMMC code (Das et al. 2015; Gorai et al. 2017a, 2017b, 2020; Sil et al.2018) to explore the fate of P-bearing species in the denser region. In the next section, we first test our model for the diffuse region to validate our results and further extend it for the more evolved stage.

4.2.1. Diffuse Cloud Model

We consider the initial elemental abundances for the diffuse cloud model as in Chantzos et al. (2020; see Table 4). The dust-to-gas ratio of 0.01 is considered in our model. We consider a photodesorption rate of 3 × 10⁻³ molecules per incident UV photon for all the molecules. Öberg et al. (2007) experimentally derived this rate from the laboratory measurement of CO ice. A nonthermal desorption constant a = 0.01 and a cosmic-ray ionization rate of 1.7 × 10⁻¹⁶ s⁻¹ are considered. We keep the gas and dust temperatures constant at 40 and 20 K, respectively.

Table 4. Initial Elemental Abundance for the Diffuse Cloud Model Considered in the CMMC Code (Chantzos et al. 2020)

Element	Abundance	Element	Abundance
H	1.0	Si+	3.2 × 10−5
He	8.5 × 10−2	Fe+	3.2 × 10−5
N	6.8 × 10−5	Na+	1.7 × 10−6
O	4.9 × 10−4	Mg+	3.9 × 10−5
C+	2.7 × 10−4	Cl+	3.2 × 10−7
S+	1.3 × 10−5	P+	2.6 × 10−7
		F+	3.6 × 10−8

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Figure 8 represents the results obtained with the CMMC code for the diffuse cloud region. The solid curve in the figure represents the case when we consider n_H = 100 cm⁻³, and the dotted curve represents the case when we consider n_H = 500 cm⁻³. Figure 8 depicts that in the ranges of A_V = 2–3 mag and n_H = 100–500 cm⁻³ we have a better agreement with the observed abundance.

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Based on the obtained results, Figure 9 shows the chemical evolution of most of the P-bearing species considered in this study. In our case, we did not have a high abundance of PH₃ under this situation. This is because of the inclusion of the destruction pathways of PH₃ by H and OH in our network. The dashed green curve shows the PH₃ when we do not consider its destruction by H and OH in both the gas and ice phases. It is interesting to note that we have a couple of orders higher abundance of PH₃ in the absence of these destruction pathways. However, these destruction pathways are essential and need to be considered before constraining the PH₃ abundance. The right panel of Figure 9 depicts that the abundance of P⁺ remains constant for this case. Figure 10 shows the final abundance of P and P⁺ with a variation of A_V . It shows that P⁺-to-P conversion is possible between A_V = 2 and 3 mag and remains invariant beyond A_V > 4 mag. For this case, we always get a higher peak abundance of PN compared to PO. At the end of the simulation, Chantzos et al. (2020) also obtained a comparatively higher abundance of PN than PO. With the Cloudy code, in Figure 4 we also have obtained PO/PN < 1. The major difference between the CMMC model or the model used by Chantzos et al. (2020) and the Cloudy code is the consideration of the physical conditions. The Cloudy code considers physical conditions more realistically. Figure 5 (diffuse cloud results with Cloudy) shows a substantial temperature variation with A_V , whereas in the case of the CMMC model we assume a constant temperature (T_gas = 40 K and T_ice = 20 K) for all A_V . In the CMMC model, a dust-to-gas ratio of ∼0.01 is considered, whereas in the diffuse cloud model using Cloudy it gives 6.594 × 10⁻³.

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4.2.2. Hot Core/Hot Corino Model

Here, we consider a three-stage starless collapsing cloud model described in Gorai et al. (2020). The first stage corresponds to an isothermal (T_dust = T_gas = 10 K) collapsing stage where density can increase from 3 × 10³ cm⁻³ to 10⁷ cm⁻³ and the visual extinction parameter (A_V ) can increase from 2 to 200. The second stage corresponds to a warm-up stage where temperature rises from 10 to 200 K, keeping A_V constant at its maximum value in 5 × 10⁴ yr. The last stage is a post-warm-up stage where density, temperature, and visual extinction are constant at their respective highest values and continue for 10⁵ yr. Based on the timescale of the initial isothermal collapsing stage, we define the hot core and hot corino case. In the hot core case, we use an isothermal collapsing timescale ∼10⁵ yr, whereas for the hot corino case a relatively longer timescale (∼10⁶ yr) is used. Thus, a total simulation timescale of 2.5 × 10⁵ yr and 1.15 × 10⁶ yr is considered for the hot core and hot corino cases, respectively. Table 5 shows the low metallic initial elemental abundance (Wakelam & Herbst 2008), which is considered here.

Table 5. Initial Elemental Abundance for the Hot Core/Corino Model Considered in the CMMC Code (Wakelam & Herbst 2008)

Element	Abundance	Element	Abundance
H2	0.5	Fe+	3.00 × 10−9
He	1.40 × 10−1	Na+	2.00 × 10−9
N	2.14 × 10−5	Mg+	7.00 × 10−9
O	1.76 × 10−4	Cl+	1.00 × 10−9
C+	7.30 × 10−5	P+	2.00 × 10−10
S+	8.00 × 10−8	F+	6.68 × 10−9
Si+	8.00 × 10−9	e−	7.31 × 10−5

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Rivilla et al. (2016) reported the observed molecular abundance with an uncertainty of ∼(0.5–5) × 10⁻¹⁰ for PO and PN. The ratio between PO and PN is in the range of 1–5. Figure 11 shows the variation of the peak abundance of PO and PN in the case of hot core and hot corino by considering various initial elemental P abundances. The peak abundances of PN and PO from our hot core (left panel) model show a good match with the observation toward the massive star-forming region when an initial elemental P abundance (P⁺ abundance) of 2 × 10⁻¹⁰–2 × 10⁻⁹ is considered.

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The solid curve of Figure 11 represents the BE of the P-bearing species as it is in KIDA. The dashed curve represents our computed BE values (BE with water c-tetramer) of phosphorous (see the BE values in Table 1) and keeps all other BE values the same as in KIDA. We notice an increase in the peak abundances of PO and PN for the inclusion of our calculated BEs for the hot corino case. We have obtained that the peak abundance of PN increases, whereas it decreases for PO for the hot core case. PO's peak abundance is always more prominent than the peak abundance (beyond the isothermal stage) of PN in both cases. Variation of the PH₃ abundance is also shown, which is far below the derived upper limit for the C-rich envelope (Agúndez et al. 2008, 2014). The changes in BE reflect a significant increase in the gas-phase abundance of PH₃ for the hot corino case, but a marginal decrease in the peak abundance of PH₃ is obtained in the case of the hot core. Additionally, in Figure 11, PH₃ abundance is shown with a solid green curve (with the BE of KIDA) when its destruction by H and OH is not included. It offers a significantly higher abundance of PH₃.

Unless otherwise stated, we always use a nonthermal desorption factor of a = 0.01. Additionally, we have tested with a reactive desorption factor of 0.023 (Nguyen et al. 2020) for PH₃, which yields roughly 2.3 times higher abundance of PH₃ in our case.

The solid lines of Figure 12 show the time evolution of the abundances of PO and PN and PO/PN for the hot core (top panel) and hot corino (bottom panel) cases. The dashed lines represent the time evolution of the abundances of PO and PN and the ratio PO/PN by considering the BEs of P-bearing species with the tetramer configuration of water. For a better understanding, we highlight the observed abundances and observed abundance ratio in the high-mass star-forming region by Rivilla et al. (2016). For the hot core case, we consider the initial P⁺ abundance ∼1.8 × 10⁻⁹, and for the hot corino case, a comparatively higher initial abundance of P⁺ (∼5.6 × 10⁻⁹) is considered to match the observational result. We have obtained an exciting difference between the PO and PN abundances during the later stages of the simulation of the post-warm-up stage. We obtain PO/PN > 1 in the late post-warm-up stage of the hot core, whereas it shows an opposite trend in the hot corino case. The higher abundance of PN in the hot corino case happens as a result of the presence of a comparatively more elevated amount of atomic nitrogen (shown in Figure 12 with the pink curve) in this case. This high abundance of atomic nitrogen undergoes PO + N →PN + O and yields PO/PN < 1 in the hot corino case. Also, the abundance of atomic P increases owing to the destruction of P-compounds. Our modeling results for hot core (with PO/PN > 1) and hot corino (with PO/PN < 1) agree well with the modeling results presented by Jiménez-Serra et al. (2018). The peak abundance of Figure 13 shows the time evolution of most of the P-bearing species with an initial P⁺ abundance of 1.8 × 10⁻⁹. It is interesting to note that at the end of the simulation timescale mainly P is locked in the form of PO and PN.

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4.2.3. Diffuse to Dense Cloud Model

To avoid any ambiguity for considering the depletion factor, here we consider that the diffuse cloud converts into a molecular cloud after a sufficient time. The adopted physical condition for this model is shown in Figure 14. It depicts that the total simulation time is divided into four steps. The first step is the diffuse stage, and it continues for the initial 10⁷ yr. The second step is the collapsing stage, whose span is for 10⁵ yr. The warm-up stage starts after the collapsing stage, and it lasts for 5 × 10⁴ yr. Finally, the post-warm-up stage continues for another 10⁵ yr. In total, the simulation timescale is 1.025 × 10⁷ yr. Initial elemental abundance for this case was taken from Chantzos et al. (2020; see Table 4). All the BEs are used from the KIDA database. Figure 15 shows the time evolution of the H, H₂, C⁺, C, and CO. It is interesting to note that during the lifetime of the diffuse cloud (∼10⁷ yr), most of the H atoms convert into H₂. Ionized carbon is converted into neutral carbon. During the collapsing stage, the conversion of C to CO takes place. Once CO has formed, it starts to deplete on the grain and forms complex organic molecules (COMs). In the warm-up stage, these CO and its related molecules desorbed back to the gas phase. Figure 16 shows the chemical evolution of the notable P-bearing species. Figure 16 shows that in the warm-up stage we have very high abundances of PN and PO. The abundance PH₃ is also significantly higher (∼10⁻¹⁰ with respect to total H nuclei) comparable with Jiménez-Serra et al. (2018). At the end of the simulation, we see that most of the P is locked in the form of PO and PN.

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4.3. An Outline of Our Modeling Results and a Comparison with the Earlier Results

Table 6 shows the peak abundances of the significant P-bearing species obtained with the two models (Cloudy and CMMC). However, due to more realistic physical conditions, the Cloudy code is preferred for the diffuse cloud region. Both the diffuse cloud models (Cloudy and CMMC) show that PN's abundance is higher than that of PO. Chantzos et al. (2020) also predicted a higher abundance of PN (4.8 × 10⁻¹¹) compared to PO (1.4 × 10⁻¹¹) in the diffuse cloud region. The abundances of PH₃ dramatically vary between the two models; this is because Cloudy does not consider the grain-surface reactions for the formation of PH₃, whereas the grain chemistry is adequately considered in CMMC. Moreover, the adopted physical condition differs between the two models. We have obtained a higher abundance of P-bearing species for our diffuse-dense model because of the consideration of the initially high abundance of P⁺ (∼2.6 × 10⁻⁷). In general, we have obtained a nice trend with the peak abundance ratio between PO and PN. We find PO/PN ratio <1 for the diffuse cloud region, ≪1 for the PDR region, and >1 for the hot core/corino (late warm-up stage) region. In the late post-warm-up stage, we have PO/PN > 1 for the hot core and <1 for the hot corino. We have noticed that the reaction PO + N → PN + O is mainly responsible for controlling this ratio.

Table 6. Peak Abundances of the P-bearing Species under Various Physical Conditions

Species	Cloudy Output		CMMC Output
	Diffuse Cloud (n/nH)	PDR (n/nH)	Diffuse Cloud (n/nH)	Hot Core $(n/{n}_{{{\rm{H}}}_{2}}$)	Hot Corino $(n/{n}_{{{\rm{H}}}_{2}}$)	Diffuse-dense ($n/{n}_{{{\rm{H}}}_{2}}$)
	(AV = 2, nH = 300 cm−3)	(AV = 5, Tgas = 50 K, Tice = 20 K	(AV = 2, Tgas = 40 K, Tice = 20 K	(initial P+ ∼ 1. 8 × 10−9)	(initial P+ ∼ 5. 6 × 10−9)	(initial P+ ∼ 2. 6 × 10−7)
		nH = 105.5 cm−3)	nH = 300 cm−3)
PN	3.33 × 10−10	4.02 × 10−09	4.6 × 10−11	1.7 × 10−10 (2.9 × 10−11)	1.4 × 10−10 (8.0 × 10−11)	1.04 × 10−7
PO	1.47 × 10−10	2.60 × 10−11	4.4 × 10−12	2.3 × 10−10 (4.9 × 10−11)	2.6 × 10−10 (6.6 × 10−11)	3.7 × 10−8
PH	1.13 × 10−10	1.82 × 10−11	9.6 × 10−12	5.5 × 10−11 (1.5 × 10−11)	6.5 × 10−11 (8.9 × 10−12)	2.3 × 10−9
PH3	6.30 × 10−31	1.16 × 10−22	8.9 × 10−15	1.9 × 10−12(2.1 × 10−10)	9.2 × 10−12 (2.3 × 10−10)	2.0 × 10−10
CP	2.03 × 10−11	3.24 × 10−12	1.8 × 10−11	3.5 × 10−12 (6.7 × 10−13)	3.1 × 10−11 (3.2 × 10−12)	1.2 × 10−10
HCP	1.38 × 10−13	3.85 × 10−14	9.1 × 10−12	9.5 × 10−12 (3.7 × 10−12)	3.5 × 10−11 (4.3 × 10−12)	1.3 × 10−9
HPO	4.43 × 10−17	1.28 × 10−19	5.9 × 10−14	1.2 × 10−11 (1.6 × 10−12)	1.5 × 10−11 (9.8 × 10−13)	3.2 × 10−9

Note. The bracketed values in the CMMC model denote the case when the destruction of PH₃ by H and OH is ignored.

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The results of our diffuse cloud modeling with the Cloudy and CMMC code are in line with those obtained by Chantzos et al. (2020). The significant difference between our diffuse cloud model with the Cloudy code and the model presented in Chantzos et al. (2020) is in consideration of realistic physical conditions instead of just considering the constant temperature. In contrast, the diffuse cloud model with the CMMC code adopted similar physical parameters to those considered in Chantzos et al. (2020). Since Cloudy was not equipped with adequate surface chemistry treatment, we did not have a notable abundance of PH₃. Our CMMC results show lower PH₃ because of its destruction by H and OH.

Charnley & Millar (1994) studied the chemistry of P in hot molecular cores. They considered that the gas-phase P-chemistry in the hot core starts from PH₃ release from the interstellar grains. In their case, PH₃ was gradually destroyed and transformed into P, PO, and PN. They noticed that the formation timescale of other P-bearing species is longer than those of the hot core. We consider the in situ formation of PH₃ via chemical reaction on interstellar grains. Like Charnley & Millar (1994), in Figure 13 we also notice that at the end of the warm-up stage the abundances of most of the P-bearing species remain in the form of P, PO, and PN.

Rivilla et al. (2016) found that PO and PN are chemically associated and formed during the cold collapse stage by the gas-phase reactions. At the end of the collapsing stage, these two species were deposited to the interstellar grain. These again desorb back to the gas phase during the warm-up stage when the temperature is around 35 K. Our CMMC model shown in Figure 12 shows the similar behavior of PO and PN. In the isothermal collapsing stage, the gas-phase abundance of these two species offers a higher abundance, while at the end of this stage these two are depleted to the grain. Once the temperature has become >35 K in the warm-up stage, it desorbed back to the gas phase.

Jiménez-Serra et al. (2018) constructed their model to explain the abundances of P-bearing species in a wide range of astrophysical conditions. They constructed a short-lived and long-lived chemical model depending on the time of the collapse. In the short-lived collapse, they stopped their calculations when the gas density reaches its maximum value. In the long-lived collapse, they considered some additional time after getting the final density. They noticed that the P and PN are the most abundant phosphorous-bearing species in the collapsing stage. Their maximum abundance is in the range (5–10) × 10⁻¹⁰. From our hot core model, in the collapsing stage (see Figure 13) we also have obtained that P and PN remain the most abundant P-bearing species, and their peak abundance varies in the range (7–27) × 10⁻¹⁰.

Jiménez-Serra et al. (2016) carried out a high-sensitivity observation toward the L1544 pre-stellar core. They were not able to identify the PN transitions. However, they predicted an upper limit of ∼4.6 × 10⁻¹³ for the abundance of PN. The first stage of Figure 13 represents the isothermal collapsing stage of a hot core. At the end of the collapsing stage, we have obtained a PN abundance of ∼5.4 × 10⁻¹⁵, which is consistent with this upper limit. Fontani et al. (2016) identified PN in various dense cores, which are in different evolutionary stages (starless, with a protostellar object, and with an ultracompact H ii region) of the intermediate- and high-mass star formation. They obtained all the transitions of PN where the temperature is <100 K, and line widths are <5 km s⁻¹, suggesting that the origin of PN is in the quiescent and cold region. Because of the lack of data of the thermal dust continuum emission (at the millimeter/submillimeter regime for all these sources), PN's abundances were not derived. Mininni et al. (2018) identified a few transitions of PN toward some of these sources (a sample of nine massive dense cores in different evolutionary stages). They had calculated the H₂ total column densities of the sources and derived the abundances of PN. For the slightly warmer region (25–30 K), Fontani et al. (2016) and Mininni et al. (2018) obtained 10⁻¹¹ and 5 × 10⁻¹², respectively. Our Figure 13 depicts that when the temperature is ∼35 K in the warm-up stage, we have a little higher PN abundance ∼2.08 × 10⁻¹¹ with respect to H₂. The main reason behind this difference is that Mininni et al. (2018) found the excitation temperature of PN to be ∼5–30 K. Since the average density (∼10^4–5 cm⁻³) of their targeted regions is below the critical density (10^5–6 cm⁻³) of the PN, they suggested a subthermal excitation of PN. The total hydrogen density can reach as high as 10⁷ cm⁻³ in our isothermal stage. In the warm-up stage, we have considered the same density. So, in our case, we have an average density, which is greater than the critical density of PN. So, a direct comparison between our model and the observation of Mininni et al. (2018) and Fontani et al. (2016) would not be appropriate. Here, we have referred to these observations to infer the enhancement of the PN abundance with the increase in temperature from 10 to 35 K only.

After releasing to the gas phase, PH₃ can be destroyed rapidly (Jiménez-Serra et al. 2018). The gas-phase formation of PH₃ can continue by the reactions R144 and R145 of Table A1 (in Appendix A). These two reactions were considered in Jiménez-Serra et al. (2018) and found to contribute marginally. Here, we have some additional destruction reactions of PH₃ by H (grain-phase reaction R4 and gas-phase reaction R149 of Table A1) and OH (grain-phase reaction R5 and gas-phase reaction R161 of Table A1), which yields a much lower PH₃ in the gas phase.

In Table 7, a summary of the PO/PN abundance ratio is listed in the different astrophysical environments, along with a comparison with the earlier literature (Tenenbaum et al. 2007; Aota & Aikawa 2012; De Beck et al. 2013; Fontani et al. 2016; Lefloch et al. 2016; Rivilla et al. 2016, 2018, 2020; Jiménez-Serra et al. 2018; Bergner et al. 2019; Chantzos et al. 2020; Bernal et al. 2021). Our modeling results agree well with the observed (Rivilla et al. 2016) and modeled (Jiménez-Serra et al. 2018; Chantzos et al. 2020) PO/PN ratios.

Table 7. Summary of the Obtained Abundance Ratio between PO and PN in Different Astrophysical Environments

References	Type of Study	PDR	Diffuse Cloud	Hot Core	Hot Corino	Shock	Comet	Other
This work	Modeling	≪1	<1	≥1 a	≥1 b , ≤ 1 c	⋯	⋯	⋯
Bernal et al. (2021)	Observation	⋯	⋯	⋯	⋯	⋯	⋯	2.7 d
Rivilla et al. (2020)	Observation	⋯	⋯	⋯	⋯	⋯	>10 e	1.7 f
Chantzos et al. (2020)	Modeling		<1
Bergner et al. (2019)	Observation	⋯	⋯	⋯	⋯	⋯	⋯	(1–3) g
Rivilla et al. (2018)	Observation	⋯	⋯	⋯	⋯	⋯	⋯	∼1.5 h
Jiménez-Serra et al. (2018)	Modeling	≪1	⋯	≥ 1	<1	≤ 1	⋯	⋯
Lefloch et al. (2016)	Observation	⋯	⋯	⋯	⋯	≈3 i	⋯	⋯
Rivilla et al. (2016)	Observation	⋯	⋯	1.8 j , 3 k	⋯	⋯	⋯	⋯
Fontani et al. (2016)	Observation	⋯	⋯	⋯	⋯	⋯	⋯	<(1.3–4.5) l
De Beck et al. (2013)	Observation	⋯	⋯	⋯	⋯	⋯	⋯	0.17–2 m
Aota & Aikawa (2012)	Modeling	⋯	⋯	⋯	⋯	∼(0.5–1.3) i , l	⋯	⋯
Tenenbaum et al. (2007)	Observation	⋯	⋯	⋯	⋯	⋯	⋯	2.2 n

Notes.

^a For the late warm-up to post-warm-up stage of the hot core model. ^b During the late warm-up to the initial post-warm-up of hot corino stage. ^c During the late post-warm-up stage. ^d Orion-KL. ^e 67P/Churyumov–Gerasimenko (Altwegg et al. 2016). ^f Average over multiple positions and velocity components toward AFGL 5142. ^g Class I low-mass protostar B1-a. ^h G + 0.693-0.03, ⁱ L1157-B1. ^j W51 e1/e2. ^k W3(OH). ^l Taking minimum and maximum values of the upper limit of beam-averaged column densities of multiple sources. ^m IK Tauri. ⁿ VY Canis Majoris.

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Here, we employ a 1D radiative transfer model (Hogerheijde & van der Tak 2000) to look into the transitions from the notable P-bearing species in the diffuse cloud region and around the more evolved stage, the hot core/corino region.

5.1. Diffuse Cloud Model

Here we model the galactic diffuse cloud region present in the line of sight toward the strong continuum source B0355+508 (Chantzos et al. 2020). We divide the entire cloud into 23 spherical shells. Since around the diffuse cloud region hardly a few monolayers of ice could survive, we consider a dust absorption coefficient with the bare grain (MRN model with no ice mantle) with a coagulation time of 10⁸ yr (see Table 1 of Ossenkopf & Henning 1994). We have used the dust emissivity (κ) by considering a power-law emissivity model,

$$\begin{eqnarray}&&\kappa ={\kappa }_{0}{(\nu /{\nu }_{0})}^{\beta },\end{eqnarray} \tag{ 3 }$$

where κ₀ = 1520 cm² gm⁻¹ and ν₀ = 1.62 × 10¹³ Hz (Ossenkopf & Henning 1994). The power-law index β is considered to be ∼1.0. We have run several cases varying the number density of the collision partners (atomic and molecular hydrogen), the kinetic temperature of the gas, the abundance profile, the Doppler broadening parameter, the size of the cloud, etc. The diffuse cloud model discussed in the earlier section shows that at the end of the simulation (∼10⁷ yr) a significant portion of hydrogen would convert into H₂. Keeping this in mind, we consider n_H = 100 cm⁻³ and ${n}_{{{\rm{H}}}_{2}}=400\,{\mathrm{cm}}^{-3}$ for this simulation.

In Figure 17, we show the absorption profile of the 1 → 0 transition of HNC as obtained from the 1D-RATRAN model. Here, we consider the different sizes of the cloud to represent the absorption observed by Chantzos et al. (2020). We notice that with the increase in the size of the cloud, the absorption is getting stronger. For the 65 pc size, we obtain a good match with the observed intensity. We use an abundance of 2 × 10⁻¹⁰ and the Doppler parameter of ∼0.2 km s⁻¹ for HNC. In Figure 18, the extracted observed spectra of HNC (1−0), CN (1−0), C³⁴S (2−1), and ¹³CO (1−0) (Chantzos et al. 2020) are shown in black, along with the modeled spectra shown in red. It is interesting to note that with the chosen parameter we can successfully explain the absorption features of HNC, CN, and C³⁴S and the emission feature of ¹³CO.

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In contrast to the modeling result of Chantzos et al. (2020) and the modeling results discussed in the earlier sections above, a combination of higher gas temperature (∼70 K) and density is required to explain the observed line profiles. We consider the gas and dust temperature to be equal in our calculation. Here, we use the collisional data file in the LAMDA database for HNC, CN, and ¹³CO with H₂. No collisional rates with H were available. For simplicity, we consider it to be the same as it is with H₂. We further have tested using the scaled rate for H (Schöier et al. 2005); no significant difference in synthetic spectra is observed. Since the collisional data file of C³⁴S was not available, we consider it similar to the C³²S from the LAMDA database. We obtain the best-fitted Doppler broadening parameters of 0.2, 0.4, 0.25, and 0.592 km s⁻¹ for HNC, CN, C³⁴S, and ¹³CO, respectively, which are consistent with the FWHM values obtained by Chantzos et al. (2020). It is to be noted that our results can simultaneously explain the absorption features of HNC, CN, and C³⁴S and the emission feature of ¹³CO.

Here, we convolve the synthetic spectra with the perfect beam size obtained for IRAM 30 m observation. We obtain a best-fitted abundance of ∼2 × 10⁻¹⁰ for HNC. Chantzos et al. (2020) obtained an abundance of HNC ∼2.1 × 10⁻¹⁰ from the modeling. Chantzos et al. (2020) calculated a ³²S/³⁴S isotopic ratio for the −17 km s⁻¹ cloud of 18.7 and a ¹²CO/¹³CO isotopic ratio of 16.7. Considering this fractionation ratio, the abundance ratio predicted by Chantzos et al. (2020) for HNC:CN:C³⁴S:¹³CO is 1: 12.5: 0.24: 576.35. Maintaining the same abundance ratio of Chantzos et al. (2020), we have 2.5 × 10⁻⁹, 4.8 × 10⁻¹¹, and 1.15 × 10⁻⁷ for CN, C³⁴S, and ¹³CO, respectively. The modeled spectra with these abundances are shown in Figure 18 with the solid red lines. In the case of C³⁴S, we have an excellent match with this abundance. But for CN and ¹³CO, we have strong absorption and emission, respectively, with these abundances. For CN and ¹³CO, an abundance of 7 × 10⁻¹⁰ and 2.4 × 10⁻⁸, respectively, shows an excellent match (see the green dashed lines).

In Figure 19, the line profiles of the P-bearing species PN, HCP, and PO in the diffuse cloud region are shown. Using the same diffuse cloud model and considering abundances obtained from the CMMC model, i.e., 4.6 × 10⁻¹¹ for PN, 9.1 × 10⁻¹² for HCP, and 4.4 × 10⁻¹² for PO, we generate the synthetic spectra for these three P-bearing molecules. For this purpose, we use the Doppler parameter to be 0.3 km s⁻¹ for PN, HCP, and PO, respectively, from Chantzos et al. (2020), which is consistent with the low FWHM values observed for other molecules in the diffuse cloud. The collisional data for HCP and PN are taken from the Basecol ¹² database. The collisional rate of PO was taken from the LAMDA ¹³ database. The collisional rate with the H atom is considered to be the same as it is with H₂. For the IRAM 30 m telescope, the beam convolution is considered using the beam sizes mentioned in Appendix B (see Tables B1–B4). Interestingly, for PN and PO we have obtained the spectra in absorption, whereas for the HCP we have obtained it in emission. To further check the excitation condition for the molecules observed and the phosphorus-bearing molecules targeted, we employ the RADEX code (van der Tak et al. 2007). For the RADEX calculation, we consider column densities for HNC, CN, C³⁴S, and ¹³CO of 0.69 × 10¹², 0.87 × 10¹³, 1.64 × 10¹¹, and 3.98 × 10¹⁴ cm⁻² and FWHM of 0.73, 0.54, 0.43, 0.82 km s⁻¹, respectively, from Table 4 of Chantzos et al. (2020) for the diffuse cloud at an offset velocity ≃ −17 km s⁻¹ with respect to the source. We consider the input kinetic temperature 40 K and number density of collision partner H as 300 cm⁻³ and a very low H₂ density 10 cm⁻³. A standard background of 2.73 K is considered. In Figures 18 and 19, the line excitation obtained with RADEX is shown with the maroon vertical line at 0 km s⁻¹. Interestingly, in Figure 18 the excitation for HNC, CN, and C³⁴S (for which absorption was observed) is very weak for the previously observed molecules. In contrast, the excitation with RADEX completely matches with the observed spectra for ¹³CO, observed in emission. Under the optically thin approximation and considering a very crude assumption, for the two-level system the critical density can be represented by the ratio between the Einstein A-coefficient in s⁻¹ and collisional rates in cm³ s⁻¹ (Shirley 2015). In Table 8, we have noted the Einstein A-coefficient and collisional rate and critical density of these transitions. It is clear from the table that, as expected, we have obtained the transitions of HCP and ¹³CO in emission (critical density low), and the rest are in absorption (critical density high) with the RATRAN code (see Figures 18 and 19).

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Table 8. Critical Density of Some Transitions under the Optically Thin Approximation

Species	Transitions	Frequency	Einstein Coefficient	Collision Rate	Critical Density
		(GHz)	(s−1)	at ∼ 10 K (cm3 s−1)	(cm−3)
CN	N = 1 − 0, J = 1/2 − 1/2, F = 3/2 − 1/2	113.169	1.182 × 10−5	9.10 × 10−12	1.299 × 106
HNC	J = 1–0	90.664	2.69 × 10−5	9.71 × 10−11	2.770 × 105
C34S	J = 2–1	96.413	1.60 × 10−5	5.06 × 10−11	3.162 × 105
13CO	J = 1–0	110.201	6.294 × 10−8	3.302 × 10−11	1.906 × 103
HCP	J = 2–1	79.903	3.61 × 10−7	6.884 × 10−11	5.244 × 103
HCP	J = 3–2	119.854	1.31 × 10−6	7.33 × 10−11	1.786 × 104
PN	J = 2–1	93.979	2.92 × 10−5	4.538 × 10−11	6.534 × 105
PN	J = 3–2	140.968	1.05 × 10−4	5.218 × 10−11	2.012 × 106
PO	J = 5/2–3/2, Ω = 1/2, F = 3–2, e	108.998	2.132 × 10−5	5.10 × 10−12	4.179 × 106
PO	J = 5/2–3/2, Ω = 1/2, F = 2–1, e	109.045	1.92 × 10−5	1.93 × 10−11	9.952 × 105
PO	J = 5/2–3/2, Ω = 1/2, F = 3–2, f	109.206	2.143 × 10−5	1.161 × 10−10	1.846 × 105
PO	J = 5/2–3/2, Ω = 1/2, F = 2–1, f	109.281	1.93 × 10−5	6.30 × 10−12	3.062 × 106

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In Figure 19, line excitations of the P-bearing molecules are shown. The excitations obtained from RADEX are well matched for HCP, which is in emission. For PN, we have seen emissions. For PO, we have obtained a strong emission with RADEX, whereas the results obtained with the RATRAN code show these in absorption.

5.2. Hot Core/Corino Model

Phosphene is an important reservoir of interstellar phosphorous (Charnley & Millar 1994). Since it is of particular interest to astronomers, the IR vibrational spectra analysis of PH₃ could be helpful to the community. The vibrational spectra of PH₃ would be beneficial for future astronomical observations with the James Webb Space Telescope. To precisely estimate the frequency and interpret the intensities, it is necessary to go beyond the harmonic approximation. The anharmonic calculations not only show a significant deviation from the harmonic calculations. Another advantage of the anharmonic analysis is that overtones and combination bands can also be analyzed with this. When anharmonicity is considered, their intensity vanishes at the harmonic level. The density functional theory (DFT) is employed to investigate the IR feature of PH₃. Within the DFT approach, the standard B3LYP functional (Becke 1993) has been used in conjunction with the 6–311G(d,p) basis set. All DFT computations have been performed employing the Gaussian suite of programs for quantum chemistry (Frisch et al. 2013). Figure 20 shows the calculated infrared feature of the ice-phase PH₃. To mimic the ice features, the PH₃ molecule is embedded in a continuum solvation field to represent local effects. The integral equation formalism (IEF) variant of the polarizable continuum model (PCM) as a default self-consistent reaction field (SCRF) method is employed with water as a solvent (Cancès et al. 1997; Tomasi et al. 2005). The implicit solvation model places the molecule of interest inside a cavity in a continuous homogeneous dielectric medium representing the solvent. The fundamental modes of vibration, along with the overtones and combination bands, are shown in Figure 20. Table 9 also notes down the wavenumbers (in cm⁻¹) and corresponding absorption coefficients (IR intensities in cm molecule⁻¹) of the fundamental bands, overtones, and combination bands.

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A comparison between our computed spectra and those obtained experimentally by Turner et al. (2015) is shown in Table 9. The computed vibrational frequencies are often scaled to resemble the experimental results. This scaling factor varies with the choice of the basis sets and implemented method. The NIST database ¹⁶ noted some of the scaling factors, which are helpful to compare with the experimental values. Based on our method and the chosen basis set, we use a vibrational scaling factor of ∼0.967. Puzzarini et al. (2014) demonstrated that the best state-of-the-art theoretical estimates have an accuracy of about 5–10 cm⁻¹ and 10–20 cm⁻¹ for fundamental and nonfundamental transitions, respectively. This accuracy allows for reliable simulations of IR spectra supporting astronomical observations. The harmonic frequencies presented in Table 9 agree with the experimental values noted in Turner et al. (2015) within a limit of about 1–15 cm⁻¹ after scaling. Our calculated values are in excellent agreement with the experimental values for the overtones and combination bands even without scaling.

Figure 21 shows the IR feature of PH₃ in the presence of various impurities. H₂O molecules would cover a significant portion of the interstellar ice mantles in the dense interstellar region. Some other species, such as CO, CO₂, CH₄, NH₃, etc. (Boogert et al. 2015; Gorai et al. 2020), can also constitute a sizable portion of the grain mantle. These molecules can influence the band profile of PH₃. We notice that the stretching of PH₃ around 2400 cm⁻¹ is getting more robust in the presence of CO₂. H₂S and SO₂ are among the main components of Venusian atmospheres (Greaves et al. 2021,Sousa-Silva et al. 2020). It is important to note that a fundamental mode of SO₂ coincides with the bending-scissoring modes of PH₃ around ∼1000–1100 cm⁻¹, which can blend the PH₃ transitions. Since H₂O is the major component of the interstellar dense cloud region, we check the effect of increasing concentration of H₂O on PH₃ separately in Figure 22. We notice that with increasing H₂O concentration, PH₃ fundamental bands are highly affected.

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In this paper, we have computed the abundance of the P-bearing species under various interstellar circumstances. We draw the following major conclusions:

• 1.  
  From this work, we found that the abundance of PH₃ is low in the diffuse cloud region and PDR. However, in the hot core region, a detectable amount of PH₃ could be produced.
• 2.  
  There appears to be a clear trend between the abundances of PO and PN. The PO/PN ratio is mainly governed by the reaction PO + N → PN + O. Since we were having a high abundance of atomic N in the diffuse clouds and PDRs (∼6.8 × 10⁻⁵) compared to hot cores (∼10⁻⁶), we found PO/PN < 1 for the diffuse cloud region and ≪1 for the PDR region. For the late warm-up to the post-warm-up region of the hot core, we obtained PO/PN ≥1. For the hot corino case, we noticed PO/PN ≥1 for the late warm-up to the initial post-warm-up phase, whereas it is ≤ 1 for the late post-warm-up stage.
• 3.  
  The BE of PH₃ with water is found to be lower than that of NH₃. In Das et al. (2018), the BE of NH₃ with a c-hexamer configuration of water was found to be 5163 K. In the current paper, the same for PH₃ is found to be only 2395 K.
• 4.  
  The abundance of PH₃ could be significantly affected by the destruction of H and OH. Without the inclusion of these pathways, any inference on the abundance of PH₃ in astrophysical environments would be inaccurate.
• 5.  
  The radiative transfer model for the diffuse cloud region can successfully explain the observed line profiles of HNC, CN, and C³⁴S in absorption and one transition of ¹³CO in emission. Some of the line profiles of the P-bearing species are estimated and proposed for future observation in the hot core region. An inverse P Cygni profile of PH₃ is expected in the hot core region.
• 6.  
  A good agreement between the calculated IR wavenumbers of PH₃ and the experimental feature of Turner et al. (2015) was seen.
• 7.  
  The stretching and bending-scissoring modes of PH₃ could be affected by the CO₂ and SO₂, respectively, in the ice.
• 8.  
  PH₃ fundamental bands are highly affected with increasing H₂O concentration in the gas phase.

M.S. gratefully acknowledges DST, the Government of India for providing financial assistance through DST-INSPIRE Fellowship [IF160109] scheme. S.S. acknowledges UGC, New Delhi, India, for providing a fellowship. B.B. acknowledges DST-INSPIRE Fellowship [IF170046] for providing financial support. S.K.M. acknowledges CSIR fellowship (Ref no. 18/06/2017(i) EU-V). P.G. acknowledges support from a Chalmers Cosmic Origins postdoctoral fellowship. A.P. acknowledges financial support from DST SERB EMR grant, 2017 (SERB-EMR/2016/005266), Banaras Hindu University, Varanasi, and The Inter-University Centre for Astronomy and Astrophysics, Pune, for associateship. A.D. acknowledges ISRO respond project (grant No. ISRO/RES/2/402/16-17) for partial financial support. This work was supported by the Japan-India Science Cooperative Program between JSPS and DST, grant No. JPJSBP120207703. This research was possible due to a Grant-In-Aid from the Higher Education Department of the Government of West Bengal.

Software: Gaussian 09 (Frisch et al. 2013), Cloudy 17.02 (Ferland et al. 2017), RATRAN (Hogerheijde & van der Tak 2000), ATRAN (Lord 1992), RADEX (van der Tak et al.2007).

All the gas- and ice-phase pathways considered in this work are shown in Table A1 with proper references.

Table A1. Reaction Pathways

Reaction	Reactions	Rate Coefficient			References
Number (Type)		α	β	γ
Gas-phase Pathways
R1 (IN)	C+ + P → P+ + C	1.0 × 10−09	0.0	0.0	McElroy et al. (2013)
R2 (IN)	H+ + P → P+ + H	1.0 × 10−09	0.0	0.0	McElroy et al. (2013)
R3 (IN)	He+ + P → P+ + He	1.0 × 10−09	0.0	0.0	McElroy et al. (2013)
R4 (IN)	NH3 + P+ → P + ${\mathrm{NH}}_{3}^{+}$	3.08 × 10−10	−0.5	0.0	McElroy et al. (2013)
R5 (IN)	Si + P+ → P + Si+	1.0 × 10−09	0.0	0.0	McElroy et al. (2013)
R6 (CR)	P + CRPHOT → P+ + e−	1.30 × 10−17	0.0	750	McElroy et al. (2013)
R7 (PH)	P + hν → P+ + e−	1.0 × 10−09	0.0	2.7	McElroy et al. (2013)
R8 (ER)	P+ + e− → P + hν	3.41 × 10−12	−0.65	0.0	McElroy et al. (2013)
R9 (NR)	P + CN → PN + C	3.0 × 10−10	⋯	⋯	Jiménez-Serra et al. (2018)
R10 (NR)	N + PH → PN + H	5.0 × 10−11	0.0	0.0	Smith et al. (2004)
R11 (NR)	N + PO → PN + O	3.0 × 10−11	−0.6	0.0	Smith et al. (2004)
R12 (NR)	N + PO → P + NO	2.55 × 10−12	0.0	0.0	Millar et al. (1987)
R13 (NR)	N + CP → PN + C	3.0 × 10−10	⋯	⋯	Jiménez-Serra et al. (2018)
R14 (NR)	C + PN+ → PN + C+	1.0 × 10−09	0.0	0.0	McElroy et al. (2013)
R15 (DR)	PN+ + e− → P + N	1.8 × 10−7	−0.5	0.0	McElroy et al. (2013)
R16 (DR)	HPN+ + e− → PN + H	1.0 × 10−7	−0.5	0.0	Millar (1991)
R17 (DR)	HPN+ + e− → P + NH	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R18 (DR)	${\mathrm{PNH}}_{2}^{+}$ + e− → PN + H2	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R19 (DR)	${\mathrm{PNH}}_{3}^{+}$ + e− → PN + H2	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R20 (CR)	PN + CRPHOT → P + N	1.30 × 10−17	0.0	250	McElroy et al. (2013)
R21 (IN)	H+ + PN → PN+ + H	1.0 × 10−9	−0.5	0.0	Millar (1991)
R22 (IN)	He+ + PN → P+ + N + He	1.0 × 10−9	−0.5	0.0	Millar (1991)
R23 (IN)	${{\rm{H}}}_{3}^{+}$ + PN → HPN+ + H2	1.0 × 10−9	−0.5	0.0	Millar (1991)
R24 (IN)	H3O+ + PN → HPN+ + H2O	1.0 × 10−9	−0.5	0.0	Millar (1991)
R25 (IN)	HCO+ + PN → HPN+ + CO	1.0 × 10−9	−0.5	0.0	Millar (1991)
R26 (NR)	N + PN → P + N2	1.0 × 10−18	0.0	0.0	Millar et al. (1987)
R27 (PH)	PN + hν → P + N	5.0 × 10−12	0.0	3.0	McElroy et al. (2013)
R28 (CR)	HPO + CRPHOT → PO + H	1.30 × 10−17	0.0	750	McElroy et al. (2013)
R29 (DR)	HPO+ + e− → PO + H	1.50 × 10−7	−0.5	0.0	Thorne et al. (1984)
R30 (DR)	HPO+ + e− → O + PH	1.50 × 10−7	−0.5	0.0	Millar (1991)
R31 (DR)	HPO+ + e− → P + O + H	1.00 × 10−7	−0.5	0.0	McElroy et al. (2013)
R32 (DR)	PO+ + e− → P + O	1.80 × 10−7	−0.5	0.0	McElroy et al. (2013)
R33 (IN)	H2O + HPO+ → PO + H3O+	3.40 × 10−10	−0.5	0.0	McElroy et al. (2013)
R34 (IN)	H2O + P+ → PO+ + H2	5.50 × 10−11	−0.5	0.0	McElroy et al. (2013)
R35 (NR)	O + HPO → PO + OH	3.80 × 10−11	0.0	0.0	Smith et al. (2004)
R36 (NR)	O + PH2 → PO + H2	4.0 × 10−11	0.0	0.0	Millar et al. (1987)
R37 (NR)	O + PH2 → H + HPO	8.0 × 10−11	0.0	0.0	Smith et al. (2004)
R38 (NR)	O + PH → PO + H	1.0 × 10−10	0.0	0.0	Smith et al. (2004)
R39 (NR)	P + OH → PO + H	6.1 × 10−11	−0.23	14.9	Jiménez-Serra et al. (2018)
R40 (NN)	O2 + P → O + PO	1.0 × 10−13	0.0	0.0	McElroy et al. (2013)
R41 (NN)	O2 + ${\mathrm{PH}}_{2}^{+}$ → PO+ + H2O	7.8 × 10−11	0.0	0.0	McElroy et al. (2013)
R42 (IN)	P+ + CO2 → PO+ + CO	4.60 × 10−10	0.0	0.0	McElroy et al. (2013)
R43 (IN)	P+ + O2 → PO+ + O	5.60 × 10−10	0.0	0.0	McElroy et al. (2013)
R44 (IN)	P+ + OCS → PO+ + CS	4.18 × 10−10	0.0	0.0	McElroy et al. (2013)
R45 (PH)	HPO + hν → PO + H	1.70 × 10−10	0.0	5.3	McElroy et al. (2013)
R46 (CR)	PO + CRPHOT → P + O	1.30 × 10−17	0.0	250	McElroy et al. (2013)
R47 (IN)	C+ + PO → PO+ + C	1.0 × 10−9	−0.5	0.0	Thorne et al. (1984)
R48 (IN)	H+ + PO → PO+ + H	1.0 × 10−9	−0.5	0.0	Thorne et al. (1984)
R49 (IN)	${{\rm{H}}}_{3}^{+}$ + PO → HPO+ + H2	1.0 × 10−9	−0.5	0.0	Thorne et al. (1984)
R50 (IN)	HCO+ + PO → HPO+ + CO	1.0 × 10−9	−0.5	0.0	Thorne et al. (1984)
R51 (IN)	He+ + PO → P+ + O + He	1.0 × 10−9	−0.5	0.0	Thorne et al. (1984)
R52 (PH)	PO + hν → P + O	3.0 × 10−10	0.0	2.0	McElroy et al. (2013)
R53 (DR)	${\mathrm{PCH}}_{2}^{+}$ + e− → HCP + H	1.50 × 10−7	−0.5	0.0	Millar (1991)
R54 (CR)	HCP + CRPHOT → CP + H	1.30 × 10−17	0.0	750	McElroy et al. (2013)
R55 (IN)	C+ + HCP → CCP+ + H	5.0 × 10−10	0.0	0.0	Millar (1991)
R56 (IN)	C+ + HCP → HCP+ + C	5.0 × 10−10	0.0	0.0	Millar (1991)
R57 (IN)	C + HCP+ → CCP+ + H	2.0 × 10−10	0.0	0.0	McElroy et al. (2013)
R58 (IN)	H+ + HCP → HCP+ + H	1.0 × 10−9	0.0	0.0	Millar (1991)
R59 (IN)	${{\rm{H}}}_{3}^{+}$ + HCP → ${\mathrm{PCH}}_{2}^{+}$ + H2	1.0 × 10−9	0.0	0.0	Millar (1991)
R60 (IN)	H3O+ + HCP → ${\mathrm{PCH}}_{2}^{+}$ + H2O	1.0 × 10−9	0.0	0.0	Millar (1991)
R61 (IN)	HCO+ + HCP → ${\mathrm{PCH}}_{2}^{+}$ + CO	1.0 × 10−9	0.0	0.0	Adams et al. (1990)
R62 (IN)	He+ + HCP → P+ + CH + He	5.0 × 10−10	0.0	0.0	Millar (1991)
R63 (IN)	He+ + HCP → PH + C+ + He	5.0 × 10−10	0.0	0.0	Millar (1991)
R64 (NR)	O + HCP → PH + CO	3.61 × 10−13	2.1	3080.0	Smith et al. (2004)
R65 (PH)	HCP + hν → CP + H	5.48 × 10−10	0.0	2.0	McElroy et al. (2013)
R66 (NR)	CCP + O → CO + CP	6.0 × 10−12	0.0	0.0	Smith et al. (2004)
R67 (IN)	CCP + He+ → He + CP + C+	5.0 × 10−10	−0.5	0.0	Millar (1991)
R68 (IN)	CCP + He+ → He + C2 + P+	5.0 × 10−10	−0.5	0.0	Millar (1991)
R69 (DR)	${\mathrm{PCH}}_{4}^{+}$ + e− → CP + C3H	7.5 × 10−8	−0.5	0.0	McElroy et al. (2013)
R70 (DR)	CCP+ + e− → C + CP	1.5 × 10−7	−0.5	0.0	Millar (1991)
R71 (DR)	CCP+ + e− → P + C2	1.5 × 10−7	−0.5	0.0	Millar (1991)
R72 (DR)	${\mathrm{PCH}}_{2}^{+}$ + e− → CP + H2	1.50 × 10−7	−0.5	0.0	Millar (1991)
R73 (DR)	HCP+ + e− → CP + H	1.50 × 10−7	−0.5	0.0	Millar (1991)
R74 (DR)	HCP+ + e− → P + CH	1.50 × 10−7	−0.5	0.0	Millar (1991)
R75 (DR)	CP+ + e− → P + C	1.00 × 10−7	−0.5	0.0	McElroy et al. (2013)
R76 (NR)	C + PH → CP + H	7.50 × 10−11	0.0	0.0	Smith et al. (2004)
R77 (CR)	CP + CRPHOT → C + P	1.30 × 10−17	0.0	250	McElroy et al. (2013)
R78 (IN)	C+ + CP → CP+ + C	1.0 × 10−9	0.0	0.0	Millar (1991)
R79 (IN)	H+ + CP → CP+ + H	1.0 × 10−9	0.0	0.0	Millar (1991)
R80 (IN)	${{\rm{H}}}_{3}^{+}$ + CP → HCP+ + H2	1.0 × 10−9	0.0	0.0	Millar (1991)
R81 (IN)	H2 + CP+ → HCP+ + H	1.0 × 10−9	0.0	0.0	McElroy et al. (2013)
R82 (IN)	H3O+ + CP → HCP+ + H2O	1.0 × 10−9	0.0	0.0	Millar (1991)
R83 (IN)	HCO+ + CP → HCP+ + CO	1.0 × 10−9	0.0	0.0	Adams et al. (1990)
R84 (IN)	He+ + CP → P+ + C + He	5.0 × 10−10	0.0	0.0	Millar (1991)
R85 (IN)	He+ + CP → P + C+ + He	5.0 × 10−10	0.0	0.0	Millar (1991)
R86 (NR)	O + CP → P + CO	4.0 × 10−11	0.0	0.0	Millar (1991)
R87 (IN)	O + CP+ → P+ + CO	2.0 × 10−10	0.0	0.0	McElroy et al. (2013)
R88 (PH)	CP + hν → C + P	1.0 × 10−9	0.0	2.8	McElroy et al. (2013)
R89 (PH)	C + P → CP + hν	1.41 × 10−18	0.03	55.0	McElroy et al. (2013)
R90 (DR)	PC2 ${{\rm{H}}}_{3}^{+}$ + e− → PH + C2H2	1.00 × 10−7	−0.50	0.0	McElroy et al. (2013)
R91 (DR)	${\mathrm{PNH}}_{2}^{+}$ + e− → PH + NH	1.50 × 10−7	−0.50	0.0	McElroy et al. (2013)
R92 (DR)	${\mathrm{PNH}}_{3}^{+}$ + e− → PH + NH2	1.50 × 10−7	−0.50	0.0	McElroy et al. (2013)
R93 (DR)	${\mathrm{PNH}}_{2}^{+}$ + e− → P + NH2	1.50 × 10−7	−0.50	0.0	McElroy et al. (2013)
R94 (DR)	${\mathrm{PNH}}_{3}^{+}$ + e− → P + NH3	3.0 × 10−7	−0.50	0.0	McElroy et al. (2013)
R95 (DR)	${\mathrm{PH}}_{2}^{+}$ + e− → PH + H	9.38 × 10−8	−0.64	0.0	McElroy et al. (2013)
R96 (DR)	${\mathrm{PH}}_{3}^{+}$ + e− → PH + H2	1.5 × 10−7	−0.5	0.0	Millar (1991)
R97 (DR)	HPN+ + e− → PH + N	1.0 × 10−7	−0.5	0.0	Millar (1991)
R98 (DR)	H2PO+ + e− → PH + OH	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R99 (DR)	H2PO+ + e− → HPO + H	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R100 (IN)	H2O + ${\mathrm{PH}}_{2}^{+}$ → PH + H3O+	1.62 × 10−10	−0.5	0.0	Adams et al. (1990)
R101 (IN)	NH3 + PH+ → P + ${\mathrm{NH}}_{4}^{+}$	3.99 × 10−10	−0.5	0.0	McElroy et al. (2013)
R102 (IN)	NH3 + PH+ → ${\mathrm{PNH}}_{2}^{+}$ + H2	5.88 × 10−10	−0.5	0.0	McElroy et al. (2013)
R103 (IN)	NH3 + PH+ →${\mathrm{PNH}}_{3}^{+}$ + H	1.11 × 10−9	−0.5	0.0	McElroy et al. (2013)
R104 (IN)	NH3 + P+ → ${\mathrm{PNH}}_{2}^{+}$ + H	2.67 × 10−10	−0.5	0.0	McElroy et al. (2013)
R105 (IN)	NH3 + ${\mathrm{PH}}_{2}^{+}$ → PH + ${\mathrm{NH}}_{4}^{+}$	3.8 × 10−10	−0.5	0.0	Adams et al. (1990)
R106 (IN)	NH3 + ${\mathrm{PH}}_{2}^{+}$ → ${\mathrm{PNH}}_{3}^{+}$ + H2	1.62 × 10−9	−0.5	0.0	McElroy et al. (2013)
R107 (NR)	O + PH2 → PH + OH	2.0 × 10−11	0.0	0.0	Smith et al. (2004)
R108 (NR)	C + HPO → PH + CO	4.0 × 10−11	0.0	0.0	Millar (1991)
R109 (CR)	PH + CRPHOT → P + H	1.3 × 10−17	0.0	250	McElroy et al. (2013)
R110 (IN)	C+ + PH → PH+ + C	1.0 × 10−9	0.0	0.0	Millar (1991)
R111 (IN)	H+ + PH → PH+ + H	1.0 × 10−9	0.0	0.0	Millar (1991)
R112 (IN)	${{\rm{H}}}_{3}^{+}$ + P → PH+ + H2	1.0 × 10−9	0.0	0.0	Adams et al. (1990)
R113(IN)	${{\rm{H}}}_{3}^{+}$ + PH → ${\mathrm{PH}}_{2}^{+}$ + H2	2.0 × 10−9	0.0	0.0	Millar (1991)
R114 (IN)	HCO+ + P → PH+ + CO	1.0 × 10−9	0.0	0.0	Adams et al. (1990)
R115 (IN)	O + HCP+ → PH+ + CO	2.0 × 10−10	0.0	0.0	Millar (1991)
R116 (IN)	O + HPO+ → PH+ + O2	2.0 × 10−10	0.0	0.0	Thorne et al. (1984)
R117 (IN)	HCO+ + PH → ${\mathrm{PH}}_{2}^{+}$ + CO	1.0 × 10−9	0.0	0.0	Millar (1991)
R118 (IN)	He+ + PH → P+ + He + H	1.0 × 10−9	0.0	0.0	Millar (1991)
R119 (IN)	He+ + HPO → PH+ + O + He	5.0 × 10−10	−0.5	0.0	Millar (1991)
R120 (IN)	He+ + HPO → PO+ + H + He	5.0 × 10−10	−0.5	0.0	McElroy et al. (2013)
R121 (NR)	H + PH → P + H2	1.50 × 10−10	0.0	416	Kaye & Strobel (1983)
R122 (IN)	O + PH+ → PO+ + H	1.0 × 10−9	0.0	0.0	Thorne et al. (1984)
R123 (IN)	HCN + PH+ → HCNH+ + P	3.06 × 10−10	−0.5	0.0	McElroy et al. (2013)
R124 (IN)	O + PN+ → PO+ + N	2.0 × 10−10	0.0	0.0	McElroy et al. (2013)
R125 (IN)	OH + P+ → PO+ + H	5.0 × 10−10	−0.5	0.0	McElroy et al. (2013)
R126 (IN)	PH+ + O2 → PO+ + OH	5.4 × 10−10	0.0	0.0	Adams et al. (1990)
R127 (DR)	PH+ + e− → P + H	1.0 × 10−7	−0.5	0.0	Thorne et al. (1984)
R128 (PH)	PH + hν → P + H	4.0 × 10−10	0.0	1.5	McElroy et al. (2013)
R129 (DR)	${\mathrm{PH}}_{3}^{+}$ + e− → H + PH2	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R130 (DR)	${\mathrm{PH}}_{2}^{+}$ + e− → P + H2	5.36 × 10−8	−0.64	0.0	McElroy et al. (2013)
R131 (DR)	${\mathrm{PH}}_{2}^{+}$ + e− → P + H + H	5.23 × 10−7	−0.64	0.0	McElroy et al. (2013)
R132 (IN)	NH3 + ${\mathrm{PH}}_{3}^{+}$ → PH2 + ${\mathrm{NH}}_{4}^{+}$	2.3 × 10−9	−0.5	0.0	Adams et al. (1990)
R133 (DR)	${\mathrm{PH}}_{4}^{+}$ + e− → PH2 + H2	1.50 × 10−7	−0.5	0.0	Charnley & Millar (1994)
R134 (CR)	PH2 + CRPHOT → PH + H	1.3 × 10−17	0.0	750.0	McElroy et al. (2013)
R135 (IN)	H+ + PH2 → ${\mathrm{PH}}_{2}^{+}$ + H	1.0 × 10−9	0.0	0.0	Millar (1991)
R136 (NR)	H + PH2 → PH + H2	6.20 × 10−11	0.0	318	Kaye & Strobel (1983)
R137 (IN)	${{\rm{H}}}_{3}^{+}$ + PH2 → ${\mathrm{PH}}_{3}^{+}$ + H2	2.0 × 10−9	0.0	0.0	Millar (1991)
R138 (IN)	HCO+ + PH2 → ${\mathrm{PH}}_{3}^{+}$ + CO	1.0 × 10−9	0.0	0.0	Millar (1991)
R139 (IN)	He+ + PH2 → P+ + He + H2	1.0 × 10−9	0.0	0.0	Millar (1991)
R140 (IN)	${\mathrm{PH}}_{2}^{+}$ + O2 → PO+ + H2O	7.8 × 10−11	0.0	0.0	Adams et al. (1990)
R141 (IN)	H2 + P+ → ${\mathrm{PH}}_{2}^{+}$ + hν	7.5 × 10−18	−1.3	0.0	Adams et al. (1990)
R142 (PH)	PH2 + hν → ${\mathrm{PH}}_{2}^{+}$ + e−	1.73 × 10−10	0.0	2.6	McElroy et al. (2013)
R143 (PH)	PH2 + hν → PH + H	2.11 × 10−10	0.0	1.5	McElroy et al. (2013)
R144 (DR)	${\mathrm{PH}}_{4}^{+}$ + e− → PH3 + H	1.50 × 10−7	−0.5	0.0	Charnley & Millar (1994)
R145 (IN)	${\mathrm{PH}}_{4}^{+}$ + NH3 → PH3 + ${\mathrm{NH}}_{4}^{+}$	2.10 × 10−9	0.0	0.0	Thorne et al. (1983)
R146 (NR)	PH2 + H → PH3	3.7 × 10−10	0.0	340	Kaye & Strobel (1983)
R147 (RA)	H2 + PH+ → ${\mathrm{PH}}_{3}^{+}$ + hν	2.40 × 10−17	−1.4	0.0	Adams et al. (1990)
R148 (NR)	H+ + PH3 → ${\mathrm{PH}}_{3}^{+}$ + H	7.22 × 10−11	0.0	886	Sousa-Silva et al. (2020)
R149 (NR)	H + PH3 → PH2 + H2	7.22 × 10−11	0.0	886	Sousa-Silva et al. (2020)
R150 (RR)	NH2 + PH3 → PH2 + NH3	1.50 × 10−12	0.0	928	Kaye & Strobel (1983)
R151 (IN)	H+ + PH3 → ${\mathrm{PH}}_{3}^{+}$ + H	2.00 × 10−9	0.0	0.0	Charnley & Millar (1994)
R152 (IN)	He+ + PH3 → ${\mathrm{PH}}_{2}^{+}$ + H + He	2.00 × 10−9	0.0	0.0	Charnley & Millar (1994)
R153 (IN)	C+ + PH3 → ${\mathrm{PH}}_{3}^{+}$ + C	2.00 × 10−9	0.0	0.0	Charnley & Millar (1994)
R154 (IN)	${{\rm{H}}}_{3}^{+}$ + PH3 → ${\mathrm{PH}}_{4}^{+}$ + H2	2.00 × 10−9	0.0	0.0	Charnley & Millar (1994)
R155 (IN)	HCO+ + PH3 → ${\mathrm{PH}}_{4}^{+}$ + CO	2.00 × 10−9	0.0	0.0	Charnley & Millar (1994)
R156 (IN)	H3O+ + PH3 → ${\mathrm{PH}}_{4}^{+}$ + H2O	2.00 × 10−9	0.0	0.0	Charnley & Millar (1994)
R157 (IN)	${\mathrm{PH}}_{3}^{+}$ + PH3 → ${\mathrm{PH}}_{4}^{+}$ + PH2	1.10 × 10−9	0.0	0.0	Smith et al. (1989)
R158 (PH)	PH3 + hν → PH2 + H	9.23 × 10−10	0.0	2.1	McElroy et al. (2013) [UMIST, following NH3]
R159 (PH)	PH3 + hν → PH + H2	2.76 × 10−10	0.0	2.1	McElroy et al. (2013) [UMIST, following NH3]
R160 (PH)	PH3 + hν → ${\mathrm{PH}}_{3}^{+}$ + e−	2.80 × 10−10	0.0	3.1	McElroy et al. (2013) [UMIST, following NH3]
R161 (NR)	PH3 + OH → H2O + PH2	2.71 × 10−11	0.0	155	Sousa-Silva et al. (2020)
R162 (DR)	PC2 ${{\rm{H}}}_{2}^{+}$ + e− → CCP + H2	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R163 (DR)	PC2 ${{\rm{H}}}_{3}^{+}$ + e− → CCP + H2 + H	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R164 (DR)	PC2 ${{\rm{H}}}_{4}^{+}$ + e− → CCP + H2 + H2	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R165 (DR)	PC3H+ + e− → CCP + CH	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R166 (DR)	PC4H+ + e− → CCP + C2H	7.5 × 10−8	−0.5	0.0	McElroy et al. (2013)
R167 (CR)	CCP + CRPHOT → C2 + P	1.3 × 10−17	0.0	375.0	McElroy et al. (2013)
R168 (CR)	CCP + CRPHOT → CP + C	1.3 × 10−17	0.0	375.0	McElroy et al. (2013)
R169 (IN)	C+ + CCP → CCP+ + C	5.0 × 10−10	−0.5	0.0	McElroy et al. (2013)
R170 (IN)	C+ + CCP → CP+ + C2	5.0 × 10−10	−0.5	0.0	McElroy et al. (2013)
R171 (IN)	H+ + CCP → CCP+ + H	1.0 × 10−9	−0.5	0.0	McElroy et al. (2013)
R172 (PH)	CCP + hν → C2 + P	1.0 × 10−10	0.0	1.7	McElroy et al. (2013)
R173 (PH)	CCP + hν → CP + C	1.0 × 10−9	0.0	1.7	McElroy et al. (2013)
R174 (IN)	C+ + HPO → HPO+ + C	1.0 × 10−9	−0.5	0.0	McElroy et al. (2013)
R175 (IN)	H+ + HPO → HPO+ + H	1.0 × 10−9	−0.5	0.0	McElroy et al. (2013)
R176 (IN)	H2O + P+ → HPO+ + H	4.95 × 10−10	−0.5	0.0	McElroy et al. (2013)
R177 (IN)	H2O + PH+ → HPO+ + H2	7.44 × 10−10	−0.5	0.0	McElroy et al. (2013)
R178 (IN)	H3O+ + P → HPO+ + H2	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R179 (IN)	H2O + PH+ → H2PO+ + H	2.04 × 10−10	−0.5	0.0	McElroy et al. (2013)
R180 (IN)	H2O + PH+ → P + H3O+	2.52 × 10−10	−0.5	0.0	McElroy et al. (2013)
R181 (IN)	H2O + ${\mathrm{PH}}_{2}^{+}$ → H2PO+ + H2	3.28 × 10−10	−0.5	0.0	McElroy et al. (2013)
R182 (IN)	${{\rm{H}}}_{3}^{+}$ + HPO → H2PO+ + H2	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R183 (IN)	H3O+ + HPO → H2PO+ + H2O	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R184 (IN)	HCO+ + HPO → H2PO+ + CO	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R185 (IN)	${\mathrm{CH}}_{3}^{+}$ + P → ${\mathrm{PCH}}_{2}^{+}$ + H	1.0 × 10−9	0.0	0.0	McElroy et al. (2013)
R186 (IN)	CH4 + P+ → ${\mathrm{PCH}}_{2}^{+}$ + H2	9.6 × 10−10	0.0	0.0	McElroy et al. (2013)
R187 (IN)	H2 + HCP+ → ${\mathrm{PCH}}_{2}^{+}$ + H	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R188 (DR)	${\mathrm{PCH}}_{3}^{+}$ + e− → CP + H2 + H	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R189 (DR)	${\mathrm{PCH}}_{3}^{+}$ + e− → HCP + H2	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R190 (DR)	${\mathrm{PCH}}_{3}^{+}$ + e− → P + CH3	3.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R191 (IN)	CH4 + PH+ → ${\mathrm{PCH}}_{3}^{+}$ + H2	1.05 × 10−9	0.0	0.0	McElroy et al. (2013)
R192 (IN)	H+ + CH2PH → ${\mathrm{PCH}}_{3}^{+}$ + H	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R193 (DR)	${\mathrm{PCH}}_{4}^{+}$ + e− → CH2PH + H2 + H	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R194 (DR)	${\mathrm{PCH}}_{4}^{+}$ + e− → HCP + H2 + H	1.5 × 10−7	−0.5	0.0	McElroy et al. (2013)
R195 (DR)	${\mathrm{PCH}}_{4}^{+}$ + e− → P + CH4	3.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R196 (IN)	CH4 + PH+ → ${\mathrm{PCH}}_{4}^{+}$ + H	5.5 × 10−11	0.0	0.0	McElroy et al. (2013)
R197 (IN)	CH4 + ${\mathrm{PH}}_{2}^{+}$ → ${\mathrm{PCH}}_{4}^{+}$ + H2	1.1 × 10−9	0.0	0.0	McElroy et al. (2013)
R198 (IN)	${{\rm{H}}}_{3}^{+}$ + CH2PH → ${\mathrm{PCH}}_{4}^{+}$ + H2	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R199 (IN)	H3O+ + CH2PH → ${\mathrm{PCH}}_{4}^{+}$ + H2O	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R200 (IN)	HCO+ + CH2PH → ${\mathrm{PCH}}_{4}^{+}$ + CO	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R201 (IN)	H+ + HC2P → HC2P+ + H	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R202 (CR)	HC2P + CRPHOT → CCP + H	1.3 × 10−17	0.0	750.0	McElroy et al. (2013)
R203 (DR)	PC2 ${{\rm{H}}}_{2}^{+}$ + e− → HC2P + H	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R204 (DR)	PC2 ${{\rm{H}}}_{3}^{+}$ + e− → HC2P + H2	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R205 (DR)	PC2 ${{\rm{H}}}_{4}^{+}$ + e− → HC2P + H2 + H	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R206 (DR)	PC3H+ + e− → HC2P + C	1.0 × 10−7	−0.5	0.0	McElroy et al. (2013)
R207 (IN)	C+ + HC2P → CCP+ + CH	5.00 × 10−10	0.0	0.0	McElroy et al. (2013)
R208 (IN)	C+ + HC2P → CP+ + C2H	5.00 × 10−10	0.0	0.0	McElroy et al. (2013)
R209 (IN)	${{\rm{H}}}_{3}^{+}$ + HC2P → PC2 ${{\rm{H}}}_{2}^{+}$ + H2	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R210 (IN)	H3O+ + HC2P → PC2 ${{\rm{H}}}_{2}^{+}$ + H2O	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R211 (IN)	HCO+ + HC2P → PC2 ${{\rm{H}}}_{2}^{+}$ + CO	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R212 (IN)	He+ + HC2P → CP+ + CH + He	5.00 × 10−10	0.0	0.0	McElroy et al. (2013)
R213 (IN)	He+ + HC2P → CP + CH+ + He	5.00 × 10−10	0.0	0.0	McElroy et al. (2013)
R214 (NN)	O + HC2P → HCP + CO	4.00 × 10−11	0.0	0.0	McElroy et al. (2013)
R215 (PH)	HC2P + hν → CCP + H	5.48 × 10−10	0.0	2.0	McElroy et al. (2013)
R216 (IN)	H+ + C4P → C4P+ + H	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R217 (CR)	C4P + CRPHOT → C3P + C	1.3 × 10−17	0.0	750.0	McElroy et al. (2013)
R218 (DR)	PC4H+ + e− → C4P + H	7.50 × 10−8	−0.5	0.0	McElroy et al. (2013)
R219 (IN)	${{\rm{H}}}_{3}^{+}$ + C4P → PC4H+ + H2	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R220 (IN)	H3O+ + C4P → PC4H+ + H2O	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R221 (IN)	HCO+ + C4P → PC4H+ + CO	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R222 (IN)	He+ + C4P → CCP+ + C2 + He	5.00 × 10−10	−0.5	0.0	McElroy et al. (2013)
R223 (IN)	He+ + C4P → CCP + ${{\rm{C}}}_{2}^{+}$ + He	5.00 × 10−10	−0.5	0.0	McElroy et al. (2013)
R224 (NN)	O + C4P → C3P + CO	1.00 × 10−11	0.0	0.0	McElroy et al. (2013)
R225 (PH)	C4P + hν → C3 + CP	5.48 × 10−10	0.0	1.7	McElroy et al. (2013)
R226 (CR)	C3P + CRPHOT → CCP + C	1.3 × 10−17	0.0	750.0	McElroy et al. (2013)
R227 (DR)	C4P+ + e− → C3P + C	1.50 × 10−7	−0.5	0.0	McElroy et al. (2013)
R228 (DR)	PC3H+ + e− → C3P + H	1.00 × 10−7	−0.5	0.0	McElroy et al. (2013)
R229 (DR)	PC4H+ + e− → C3P + CH	7.50 × 10−8	−0.5	0.0	McElroy et al. (2013)
R230 (IN)	${{\rm{H}}}_{3}^{+}$ + C3P → PC3H+ + H2	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R231 (IN)	H3O+ + C3P → PC3H+ + H2O	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R232 (IN)	HCO+ + C3P → PC3H+ + CO	1.00 × 10−9	−0.5	0.0	McElroy et al. (2013)
R233 (IN)	He+ + C3P → ${{\rm{C}}}_{3}^{+}$ + P + He	5.00 × 10−10	−0.5	0.0	McElroy et al. (2013)
R234 (IN)	He+ + C3P → C3 + P+ + He	5.00 × 10−10	−0.5	0.0	McElroy et al. (2013)
R235 (NN)	O + C3P → CCP + CO	4.00 × 10−11	0.0	0.0	McElroy et al. (2013)
R236 (PH)	C3P + hν → C2 + CP	5.00 × 10−10	0.0	1.8	McElroy et al. (2013)
R237 (CR)	CH2PH + CRPHOT → HCP + H2	1.3 × 10−17	0.0	750.0	McElroy et al. (2013)
R238 (IN)	C+ + CH2PH → HC2P+ + H2	1.00 × 10−9	0.0	0.0	McElroy et al. (2013)
R239 (IN)	He+ + CH2PH → PH+ + CH2 + He	5.00 × 10−10	0.0	0.0	McElroy et al. (2013)
R240 (IN)	He+ + CH2PH → PH + ${\mathrm{CH}}_{2}^{+}$ + He	5.00 × 10−10	0.0	0.0	McElroy et al. (2013)
R241 (NN)	O + CH2PH → PH2 + CO + H	4.00 × 10−11	0.0	0.0	McElroy et al. (2013)
R242 (PH)	CH2PH + hν → CH2 + PH	9.54 × 10−10	0.0	1.8	McElroy et al. (2013)
R243 (DR)	PC2 ${{\rm{H}}}_{2}^{+}$ + e− → P + C2H2	1.00 × 10−7	−0.5	0.0	McElroy et al. (2013)
R244 (IN)	C2H2 + PH+ → PC2 ${{\rm{H}}}_{2}^{+}$ + H	1.30 × 10−9	0.0	0.0	McElroy et al. (2013)
R245 (IN)	C2H2 + ${\mathrm{PH}}_{2}^{+}$ → PC2 ${{\rm{H}}}_{2}^{+}$ + H2	1.40 × 10−9	0.0	0.0	McElroy et al. (2013)
R246 (IN)	C2H2 + ${\mathrm{PH}}_{3}^{+}$ → PC2 ${{\rm{H}}}_{3}^{+}$ + H2	5.80 × 10−10	0.0	0.0	McElroy et al. (2013)
Ice-phase/Grain-surface Pathways
R1	gH + gP → gPH	⋯	⋯	⋯	Chantzos et al. (2020)
R2	gH + gPH → gPH2	⋯	⋯	⋯	Chantzos et al. (2020)
R3	gH + gPH2 → gPH3	⋯	⋯	⋯	Chantzos et al. (2020)
R4	gH + gPH3 → gPH2 + gH2	⋯	⋯	⋯	This work, Sousa-Silva et al. (2020)
R5	gOH + gPH3 → gPH2 + gH2O	⋯	⋯	⋯	This work, Sousa-Silva et al. (2020)
R6	gPH3 → PH3	⋯	⋯	⋯	Chantzos et al. (2020)

Note. CR refers to cosmic rays, IN to ion−neutral reactions, NR to neutral−radical reactions, NN to neutral−neutral reactions, RR to radical–radical reactions, RA to radiative association reactions, ER to electronic recombination reactions for atomic ions, DR to dissociative recombination reactions for molecular ions, PH to photodissociation reactions, and hν to a photon. Grain-surface species are denoted by the letter "g."

Download table as:  ASCIITypeset images: 1 2 3 4 5

Here, we note the intensities of various P-bearing species in the diffuse and hot core regions. We also highlight the observability of these transitions with IRAM 30 m, Herschel (not operational), and SOFIA. We use the ATRAN program ¹⁷ to check the atmospheric transmission around these frequencies for the ground-based and space-based observations.