A long-term goal of exoplanet studies is the identification and detection of biosignature gases. Beyond the most discussed biosignature gas O 2 , only a handful of gases have been considered in detail. In this study, we evaluate phosphine (PH 3 ). On Earth, PH 3 is associated with anaerobic ecosystems, and as such, it is a potential biosignature gas in anoxic exoplanets. We simulate the atmospheres of habitable terrestrial planets with CO 2 - and H 2 -dominated atmospheres and find that PH 3 can accumulate to detectable concentrations on planets with surface production fluxes of 10 10 to 10 14 cm −2 s −1 (corresponding to surface concentrations of 10s of ppb to 100s of ppm), depending on atmospheric composition and ultraviolet (UV) irradiation. While high, the surface flux values are comparable to the global terrestrial production rate of methane or CH 4 (10 11 cm −2 s −1 ) and below the maximum local terrestrial PH 3 production rate (10 14 cm −2 s −1 ). As with other gases, PH 3 can more readily accumulate on low-UV planets, for example, planets orbiting quiet M dwarfs or with a photochemically generated UV shield. PH 3 has three strong spectral features such that in any atmosphere scenario one of the three will be unique compared with other dominant spectroscopic molecules. Phosphine's weakness as a biosignature gas is its high reactivity, requiring high outgassing rates for detectability. We calculate that tens of hours of JWST (James Webb Space Telescope) time are required for a potential detection of PH 3 . Yet, because PH 3 is spectrally active in the same wavelength regions as other atmospherically important molecules (such as H 2 O and CH 4 ), searches for PH 3 can be carried out at no additional observational cost to searches for other molecular species relevant to characterizing exoplanet habitability. Phosphine is a promising biosignature gas, as it has no known abiotic false positives on terrestrial planets from any source that could generate the high fluxes required for detection.

1. Introduction

Life makes use of thousands of volatile molecular species that could contribute toward a biosphere and its associated atmospheric spectrum. Some of these volatiles may accumulate in a planetary atmosphere and be remotely detectable; these are commonly called “biosignature gases.”

Theoretical studies of biosignature gases have been recently heavily reviewed elsewhere (Seager et al., 2016; Grenfell, 2018; Kiang et al., 2018; Schwieterman et al., 2018).

Prominent biosignature gases on Earth are those that are both relatively abundant and spectroscopically active (primarily O 2 , its photochemical by-product O 3 , and also CH 4 and N 2 O). Other gases that are not prominent in Earth's atmosphere but might be prominent in exoplanet atmospheres have also been studied, for example, dimethyl sulfide, dimethyl disulfide, and CH 3 Cl (Pilcher, 2003; Segura et al., 2005; Domagal-Goldman et al., 2011).

The next generation telescopes will open the era of the study of rocky exoplanet atmospheres. The James Webb Space Telescope (JWST, planned for launch in 2021) is the most capable for transmission spectra studies of a handful of the most suitable rocky planets transiting bright M dwarf stars (Gardner et al., 2006), whereas ESAs Atmospheric Remote-sensing Infrared Exoplanet Large-survey (ARIEL, planned for launch in 2028) may be able to detect atmospheric components on super-Earths around the smallest M dwarf stars (Pascale et al., 2018). Large ground-based telescopes now under construction, that is, Giant Magellan Telescope, Extremely Large Telescope, and Thirty Meter Telescope (Johns et al., 2012; Tamai and Spyromilio, 2014; Skidmore et al., 2015), can also reach M dwarf star rocky planets by direct imaging, with the right instrumentation.

To the best of our knowledge, phosphine (PH 3 ) has not yet been evaluated as a biosignature gas. In Earth's atmosphere, PH 3 is a trace gas. It is possible, however, that biospheres on other planets could accumulate significant detectable PH 3 levels. In particular, anoxic biospheres where life would not be heavily dependent on oxygen could produce PH 3 in significantly higher quantities than on modern Earth (Bains et al., 2019b).

Astronomical observations find that phosphine is spectroscopically active and present in stellar atmospheres (namely carbon stars) and in the giant planet atmospheres of Jupiter and Saturn (Bregman et al., 1975; Tarrago et al., 1992; Agúndez et al., 2014). In T dwarfs and giant planets, PH 3 is expected to contain the entirety of the atmospheres' phosphorus in the deep atmosphere layers (Visscher et al., 2006), where it is sufficiently hot for PH 3 formation to be thermodynamically favored.

In both Jupiter and Saturn, phosphine is found on the high observable layers at abundances (4.8 and 15.9 ppm, respectively) several orders of magnitude higher than those predicted by thermodynamic equilibrium (Fletcher et al., 2009). This overabundance of PH 3 occurs because chemical equilibrium timescales are long when compared with convective timescales (Noll and Marley, 1997). PH 3 forms in the hotter deep layers of the atmosphere (temperatures ≿ 800 K) and is mixed upward, so that the PH 3 inventory at the cloud-top is replenished. In every astronomical body where phosphine has been detected thus far, other than on Earth, there are regions with high enough temperatures for PH 3 to be the thermodynamically favored phosphorus species.

It has been postulated that elemental phosphorus species originating from the photolysis of phosphine are responsible for the red coloring of Jupiter's red spot and other Jovian chromophores (Prinn and Lewis, 1975), although this hypothesis has not achieved wide community acceptance (Noy et al., 1981; Kim, 1996). For a review of chemical species that are current candidates for the chromophores of Jupiter, see Carlson et al. (2016) and references therein.

Phosphine has not been detected in the observable layers of ice giants, such as Uranus and Neptune (Burgdorf et al., 2004; Moreno et al., 2009), despite these planets having sufficiently hot layers to produce PH 3 and strong convection currents, which could transport PH 3 to observable altitudes. Observations put the P/H abundance in Uranus and Neptune at an upper limit of <0.1 solar P/H, which is significantly lower than expected (Teanby et al., 2019).

In this work, we critically assess phosphine as a biosignature gas. We first summarize in what circumstances phosphine is generated by life on Earth (Sections 2.1 and 2.2). We next review the known destruction mechanisms for PH 3 (Section 2.3) and describe our inputs and methods for the assessment of the detectability of PH 3 in a variety of planetary scenarios (Section 3). We then present our results (Section 4): here, we first calculate surface fluxes and associated atmospheric abundances required for the remote spectroscopic detection of PH 3 in transmission and emission spectra (Section 4.1). We then highlight the properties of the PH 3 spectrum that help distinguish it from other molecules (Section 4.2). Next, we present thermodynamic calculations that show that PH 3 , in temperate planets, has no substantial false positives as a biosignature gas (Section 4.3). We conclude with a discussion of our results (Section 5).

2. Phosphine Sources and Sinks

On Earth, phosphine is associated with biological production in anaerobic environments and anthropogenic production via a multitude of industrial processes. PH 3 has low mean production rates on Earth, but it is a mobile gas and is found globally, albeit in trace amounts, in the atmosphere. Below, we summarize the known emissions of PH 3 on Earth (Section 2.1), phosphine's association with life (Section 2.2), and its known destruction mechanisms (Section 2.3).

2.1. Phosphine emissions on Earth

Phosphine is a ubiquitous trace component of the atmosphere on modern Earth (Morton and Edwards, 2005). About 10% of the phosphorus in the atmosphere is PH 3 ; the major phosphorus form is phosphate, mostly as phosphoric acid (Elm et al., 2017). Although PH 3 is found everywhere in Earth's atmosphere, its atmospheric abundance is widely variant, with high concentration regions sometimes having more PH 3 than low concentration areas by a factor of 10,000 (Pasek et al., 2014).

Phosphine has been found worldwide in the lower troposphere of Earth in the ppq to ppb range in daytime, with higher nighttime concentrations than at daytime (due to inhibited ultraviolet [UV]-induced oxidation) (Gassmann, 1994; Gassmann et al., 1996; Glindemann et al., 1996b, 2003; Ji-ang et al., 1999; Han et al., 2000; Zhu et al., 2006a, 2007a, 2007b; Li et al., 2009; Hong et al., 2010a; Zhang et al., 2010). In the high troposphere, PH 3 was found at a peak of 7 ppt during daylight (Glindemann et al., 2003; Han et al., 2011b). This implies that sunlight does not lead to complete destruction of PH 3 , unlike previous suggestions (Glindemann et al., 2003; Han et al., 2011b). A sample of locally measured gaseous PH 3 concentrations in a variety of environments on Earth, ranging from ppq to ppb (ng/m3 to μg/m3), can be found in Fig. 1.

FIG. 1. Measurements of phosphine concentrations in Earth's atmosphere. Study number shown in x axis (references below) and y axis showing the span of locally measured concentration of gaseous PH 3 in units of ng/m3, with maximum values of 600.2 and 1259 ng/m3 (corresponding to concentrations ranging between ppq and ppb). Green bars: marshlands and paddy fields. Black bars: industrial environments. Red bar: Namibia (rural environment). White bars: Arctic and Antarctic environments. Yellow bars: upper troposphere. Blue bars: oceanic samples (coastal and open ocean). References for studies shown: (1) Han et al. (2011a); (2) Han et al. (2000); (3) Niu et al. (2013); (4) Glindemann et al. (1996a); (5) Zhang et al. (2010); (6) Glindemann et al. (1996a); (7) Zhu et al. (2007a, 2007b); (8) Zhang et al. (2010); (9, 10) Glindemann et al. (2003); (11) Li et al. (2009); (12) Zhu et al. (2007a, 2007b); (13) Gassmann et al. (1996); (14) Glindemann et al. (2003); (15) Geng et al. (2005), Han et al. (2011b); (16) Hong et al. (2010a). We do not include measurements of “Matrix-Bound Phosphine” (MBP), material that releases PH 3 when a matrix is treated with high temperatures and strong acid or alkali. Figure adapted from the work of Bains et al. (2019a). Please see Bains et al. (2019a) for more details on MBP and environmental PH 3 production. Gaseous PH 3 is found in multiple altitudes in Earth's atmosphere above a wide variety of environments, in concentrations ranging from ppq to ppb. MBP, Matrix-Bound Phosphine; PH 3 , phosphine. Color images are available online.

On Earth, a significant source of phosphine emissions is anthropogenic activity. Because of its broad toxicity to aerobic organisms,* PH 3 is widely used in the agricultural industry as a rodenticide and insecticide (Devai et al., 1988; Bingham, 2001; Glindemann et al., 2005; Perkins et al., 2015; Chen et al., 2017). PH 3 is also used commercially, for example, as a doping agent (Budavari et al., 1996). However, PH 3 emissions linked to biological activity are believed to form the majority of atmospheric PH 3 (Glindemann et al., 2005; Morton and Edwards, 2005). Evidence for the association of PH 3 with anaerobic biology is presented in Section 2.2.

2.2. Biological production of phosphine

All life on Earth relies on phosphorous compounds. The biological phosphorus cycle is heavily, but not exclusively, reliant on phosphates. Other, less oxidized, phosphorus-containing molecules also play a crucial role in the phosphorus cycle. The exact role of phosphine in the Earth's global phosphorus cycle is not yet fully known. It is, however, likely that, similar to other reduced phosphorus species, PH 3 also has an important role in the global cycling of this essential element. For more details on the role of phosphines in the Earth's phosphorus cycle, please see Appendix E.

Biological phosphine production is associated with microbial activity in environments that are strictly anoxic (lacking oxygen). This finding is in alignment with the fact that the toxicity of PH 3 is intrinsically linked to its interference with O 2 -dependent metabolism (Bains et al., 2019b).

The argument that phosphine is associated with anaerobic life is strengthened by its detection in a wide variety of ecosystems with anoxic niches, including above penguin colonies, rich in bird guano, where it reaches abundances of 300 ppt (Ji-ang et al., 1999; Zhu et al., 2006a, 2007a; Li et al., 2009; Hong et al., 2010a); animal intestinal tracts, flatus, and feces† (Gassmann and Glindemann, 1993; Eismann et al., 1997a; Chughtai and Pridham, 1998; Zhu et al., 2006a, 2006b, 2014); paddy fields‡ (Han et al., 2011a; Chen et al., 2017); rivers and lakes (Feng et al., 2008; Geng et al., 2010; Hong et al., 2010b; Han et al., 2011b; Ding et al., 2014); wetlands and marshlands (Devai and Delaune, 1995; Eismann et al., 1997b; Glindemann et al., 1996b); and landfills and sludges (Roels and Verstraete, 2004; Ding et al., 2005a, 2005b). Several studies have also reported the production of PH 3 from mixed bacterial cultures in the laboratory (Rutishauser and Bachofen, 1999; Jenkins et al., 2000; Liu et al., 2008); in one case, bacteria turning half of the phosphorus in the culture medium (∼180 mg/L) into PH 3 in 56 days (Devai et al., 1988).

Despite a large body of robust circumstantial evidence for the production of phosphine by life, the exact mechanisms for biologically associated production of PH 3 are still debated. Recent work postulates that PH 3 production may be associated with the microbial tricarboxylic acid cycle of Enterobacteriaceae (Fan et al., 2020), but the exact metabolic pathway leading to PH 3 production in anaerobic organisms remains unknown. However, we note that the absence of a known enzymatic mechanism is not evidence for the absence of biological production. The synthetic pathways for most of life's natural products are not known, and yet their origin is widely accepted to be biological because of the implausibility of their abiotic synthesis, their obligate association with life, and their chemical similarity to other biological products. For example, a recently published, manually curated, database of natural molecules produced by life on Earth contains ∼220,000 unique molecules of biological origin, produced by thousands of species (Petkowski et al., 2019a), whereas the number of known, experimentally elucidated, metabolic pathways from organisms belonging to all three domains of life is only ∼2720 (Caspi et al., 2017). Further examples of the complexities in discovering metabolic pathways for molecules associated with biological activity are provided in Appendix E.

There are two proposed explanations for the production of phosphine in anoxic ecosystems [reviewed in Glindemann et al. (1998), Roels and Verstraete (2001), Roels and Verstraete (2004), Roels et al. (2005), Bains et al. (2019a, 2019b)]:

(1) PH 3 is directly produced by anaerobic bacteria from environmental phosphorus. (2) PH 3 is indirectly produced by anaerobic bacteria. Anoxic fermentation of organic matter by anaerobic bacteria results in acid products; these acid products, in turn, could react with inorganic metal phosphides, like those present as trace elements in scrap metal, resulting in phosphine generation.

Of the two proposed explanations for biologically associated production of phosphine, we argue that the direct production as a result of metabolic activity of anaerobic bacteria is the most likely. Our reasoning is based on two lines of evidence:

(a) PH 3 has been detected in significant amounts in bacterial cultures in controlled laboratory experiments, where no metal phosphides were present, making the indirect acid-dependent production of PH 3 an unlikely scenario (Devai et al. , 1988; Glindemann et al. , 1996b; Jenkins et al. , 2000; Schink and Friedrich, 2000; Ding et al. , 2005a, 2005b; Liu et al. , 2008).

(b) Several independent studies found that PH 3 was detected in feces from evolutionarily distant animals, inhabiting diverse environments, for example, insects, birds, and mammals (including humans) (Gassmann and Glindemann, 1993; Chughtai and Pridham, 1998; Zhu et al., 2014). It is implausible that there is a significant presence of contaminant metal phosphides in the guts of all the animals, which would be required for an indirect acid-dependent production of PH 3 .

We end this introduction to the biological association of phosphine by noting that thermochemical studies on the feasibility of the production of PH 3 in temperate environments have found no plausible thermodynamically favored abiotic pathways, and as such, PH 3 has no substantial false positives for life (see Section 4.3 and Appendix C) (Bains et al., 2019a). Conversely, the production of PH 3 under anoxic conditions by living systems can be thermochemically favorable (Bains et al., 2019a), and biological functions that are accomplished through energy consuming reactions are not uncommon (Bains et al., 2019b). PH 3 could be used by life to perform complex functions that would warrant an energetic investment, such as signaling or a defense mechanism (Bains et al., 2019b). For more information about PH 3 in the context of terrestrial biology and the thermodynamic feasibility of PH 3 production by life, see Bains et al. (2019a, 2019b) and Appendices C–E.

2.3. Phosphine chemistry in the atmosphere

Within an atmosphere, phosphine is destroyed by the radicals O, H, and OH in reactions, which are thought to be first order with respect to its reactants and second order overall. PH 3 can also be regenerated by reaction of PH 2 with H and directly photolyzed by UV radiation. These processes are summarized below and discussed in Section 5.2.

Reaction rate constants are expressed via the Arrhenius equation:

where k is the reaction rate constant in units of cm3 s−1, A is a constant in units of cm3 s−1, E is the activation energy in units of J mol−1, R is the gas constant in units of J mol−1 K−1, and T is temperature in K.

2.3.1. Destruction by OH radicals

Oxidation with OH radicals is thought to be the main sink for phosphine in Earth's atmosphere via the reaction (Cao et al., 2000; Glindemann et al., 2005; Elm et al., 2017):

For this reaction, A = 2.71 × 10−11 cm3 s−1, E = 1.29 kJ mol−1, corresponding to k PH 3 ,OH = 2 × 10−11 cm3 s−1 at T = 288 K (Fritz et al., 1982). The lifetime of PH 3 due to OH reactions is calculated to be 28 h at night and 5 h in daytime, with the difference controlled by the concentration of UV-generated OH (Glindemann et al., 2003). The destruction of PH 3 by OH in the atmosphere eventually leads to phosphoric acid, which in turn contributes to the soluble phosphates found in rainwater (Lewis et al., 1985; Elm et al., 2017).

2.3.2. Destruction by O radicals

Phosphine also reacts very rapidly with atomic oxygen (on Earth, generated by photolyzed ozone), with reaction:

For this reaction, A = 4.75 × 10−11 cm3 s−1, E ≈ 0, corresponding to k PH 3 ,O = 5 × 10−11 cm3 s−1 at T = 288 K (i.e., temperature-independent in 208–423 K, Nava and Stief, 1989). Because atomic oxygen is less abundant than OH in Earth's atmosphere and troposphere, PH 3 destruction by OH is still the dominant route despite PH 3 reacting with O at a higher rate. On Earth, OH radicals described above are partially generated from the interaction between O radicals and water vapor, and so the reactions that produce the two radical species are not happening in isolation (Jacob, 1999). In anoxic atmospheres, however, the main source of OH and H radicals is the photolysis of water vapor (Hu et al., 2012).

2.3.3. Destruction by H radicals

Phosphine can be destroyed by the H radical via the reaction:

For this reaction, A = 7.22 × 10−11 cm3 s−1, E = 7.37 kJ mol−1, corresponding to k PH 3 ,H = 3 × 10−12 cm3 s−1 at T = 288 K (Arthur and Cooper, 1997). The reaction of PH 3 with the H radical is most relevant to H 2 -rich atmospheres (Seager et al., 2013b).

2.3.4. Recombination from H radicals

Phosphine can be regenerated by the radical recombination reaction:

With rate constant k PH 2 ,H = 3.7 × 10−10 exp(−340K/T) cm3 s−1, corresponding to k PH 2 ,H = 1.1 × 10−10 cm3 s−1 at T = 288 K (Kaye and Strobel, 1984). If [PH 2 ] is high, this reaction can be a major reformation pathway for PH 3 (see Section 5.2).

2.3.5. Destruction through UV radiation

UV radiation is thought to directly photolyze phosphine with unit quantum efficiency upon absorption of irradiation at wavelengths ≤230 nm (Visconti, 1981; Kaye and Strobel, 1984):

This photolysis reaction is not relevant on UV-shielded planets (e.g., modern Earth with its ozone layer) but could be relevant on anoxic planets where UV radiation may penetrate deeper into lower altitudes of the atmosphere.

Overall, phosphine is destroyed by UV irradiation, through both direct photolysis and reactions with UV-generated radical species. The daytime–nighttime PH 3 concentration difference on Earth is large due to the generation of radicals by UV irradiation during the day and their comparative absence at night. However, PH 3 has been detected at concentrations of up to 7 ppt (2.45 ng m−3) during daylight in Earth's high troposphere (Glindemann et al., 2003). PH 3 accumulates in the dry upper troposphere on Earth because ozone attenuation of UV and lack of OH-producing H 2 O result in low abundances of OH radicals, which slows the PH 3 destruction and its return to the surface in the form of phosphates (Frank and Rippen, 1987; Glindemann et al., 2003).

2.3.6. Solubility and aerosol formation

Phosphine does not easily stick to aerosols and has very low water solubility (Fluck, 1973). PH 3 is therefore a very mobile gas that is less likely to wash out and fall to the surface than other gases, such as hydrogen sulfide, methanethiol, and ammonia (Glindemann et al., 2003).

UV photolysis of phosphine in the presence of hydrocarbons could lead to the formation of complicated alkyl-phosphines (Guillemin et al., 1995, 1997). Atmospheres that are prone to high concentrations of hydrocarbon radicals, such as H 2 -rich atmospheres, could therefore lead to the creation of organophosphine hazes. The plausibility of organophosphine haze formation is discussed further in Section 5.2.

3. Inputs and Methods for the Assessment of Detectability

We have assessed the spectral distinguishability, atmospheric survival, and observational detectability of phosphine in anoxic exoplanets. In this section, we first describe the choice of molecular inputs used for our spectral analyses (Section 3.1). We then provide a brief outline of the photochemical method used to calculate the distribution of molecules throughout the atmosphere (Section 3.2). Finally, we outline the method and detectability criteria for the simulations of observational spectra (Section 3.3).

3.1. Molecular inputs

Molecular spectra can be represented in various forms to best serve as input for spectral representations and atmospheric models. For the comparison of phosphine with other major components of atmospheres (see Section 4.2), we have used cross sections calculated from the most complete spectra available. The PH 3 molecular cross sections come from the recently variationally computed PH 3 line list (Sousa-Silva et al., 2015) and the total internal partition function calculated in the work of Sousa-Silva et al. (2014). For all temperatures under 800 K (which includes all temperate environments), this PH 3 line list represents a complete spectrum containing over 16 billion transitions between 7.5 million energy levels. Even at low temperatures, it is recommended that complete line lists are used for spectral simulations; complete line lists allow for more representative cross sections with improved band shapes when compared with experimental or calculated spectra at room temperature. The carbon dioxide line list is from HITEMP (Rothman et al., 2010). All other molecular cross sections are simulated by using complete, theoretically calculated, line lists from the ExoMol database (Yurchenko et al., 2011; Yurchenko and Tennyson, 2014; Sousa-Silva et al., 2015; Tennyson et al., 2016).

For the calculation of the transmission and thermal emission spectra of the model atmospheres (see Section 4.1), molecular opacities for phosphine are adopted from the ExoMol database (Sousa-Silva et al., 2015; Tennyson et al., 2016). For all other molecules, we used the HITEMP and the HITRAN 2016 databases (Rothman et al., 2010; Gordon et al., 2017). Molecular cross sections in the UV region, used to calculate photolysis rates (see Section 3.2), were obtained from the absorption cross sections compendium of Ranjan and Sasselov (2017) and from the work of Chen et al. (1991) via the MPI-Mainz Spectral Atlas (Keller-Rudek et al., 2013).

3.2. Photochemical modeling

In this subsection, we provide a brief description of the photochemical model used to calculate the concentration of phosphine as a function of altitude for a range of PH 3 surface fluxes. We also describe the atmospheric and stellar scenarios considered in our photochemical model. We found that it is necessary to use a photochemical model instead of the approximation of fixed radical profiles (Seager et al., 2013b) because of the intense reactivity of PH 3 , which can drastically alter the radical profiles of an atmosphere. In particular, at high PH 3 fluxes, the radical concentrations are suppressed due to reactions with PH 3 , meaning that PH 3 can build up to much higher concentrations than a fixed radical profile approximation would predict.

3.2.1. Photochemical model

We adapt the photochemical model of Hu et al. (2012) to calculate atmospheric composition for different planetary scenarios. The model is detailed in the work of Hu et al. (2012); in brief, the code calculates the steady-state chemical composition of an exoplanetary atmosphere by solving the one-dimensional chemical transport equation. The model treats up to 800 chemical reactions, photochemical processes (i.e., UV photolysis of molecules), dry and wet deposition, surface emission, thermal escape of H and H 2 , and formation and deposition of elemental sulfur and sulfuric acid aerosols. The model is designed to have the flexibility of simulating both oxidized and reduced conditions. UV and visible radiation in the atmosphere is computed by the delta-Eddington two-stream method. The code has been validated by reproducing the atmospheres of modern Earth and Mars. The code and extensive application examples are described in several studies (Hu et al., 2012; Hu et al., 2013; Seager et al., 2013a, 2013b; Hu and Seager, 2014). In calculating convergence, we required that the chemical variation timescale of significant species (>1 cm−3) to be at least 1019 s, that is, longer than the age of the universe.

We added phosphine to the model of Hu et al. (2012). We considered surficial production as the only source of PH 3 , and rainout, photolysis, and reactions with the main radical species O, H, and OH as the sinks (see Section 2.3). We take the Henry's Law constant for PH 3 from the work of Fu et al. (2013) via Sander (2015). For photolysis, we take the PH 3 UV cross sections at 295 K from the work of Chen et al. (1991). We follow Kaye and Strobel (1984) in taking the branching ratio of this reaction to be unity and take the quantum yield of PH 3 photolysis to be qλ= 1 for λ < 230 nm and qλ= 0 for λ > 230 nm. For the reactions with OH, O, and H, we take the rate constants from the works of Fritz et al. (1982), Nava and Stief (1989), and Arthur and Cooper (1997), as detailed in Section 2.3. We are unaware of geochemical constraints on the dry deposition velocity of PH 3 ; we take this value to be 0 cm s−1, which could lead to an overestimation of PH 3 accumulation rates (see Section 5.2 for a discussion of possible PH 3 deposition). However, our approach neglects the possibility that atmospheric photochemistry may generate PH 3 . In particular, we do not consider the recombination reaction, PH 2 + H → PH 3 . This may lead to underestimating PH 3 accumulation, especially in H 2 -dominated atmospheres where H abundances are high.

3.2.2. Planetary scenarios

We model the atmospheres of Earth-sized, Earth-mass planets with two bulk atmospheric compositions: a H 2 -dominated atmosphere and a CO 2 -dominated atmosphere. We focus on H 2 -dominated atmospheres because their low mean molecular masses make them amenable to characterization via transmission spectroscopy (Batalha et al., 2015). We focus on CO 2 -dominated atmospheres as an oxidizing end-member to complement the reducing H 2 -dominated case, and because early Earth is thought to have had a CO 2 -rich atmosphere (Kasting, 1993). We only consider anoxic atmospheres because O 2 -rich atmospheres are likely to have large quantities of OH radicals, which rapidly destroy phosphine (see Section 2.3). Additionally, the aerobic metabolism of O 2 -dependent life is likely to be incompatible with widespread PH 3 biological production (Bains et al., 2019b).

Our atmospheres correspond to the H 2 - and CO 2 -dominated benchmark scenarios of Hu et al. (2012), with the key difference that we do not set the rainout rates of H 2 , CO, CH 4 , C 2 H 6 , or O 2 to zero, as Hu et al. (2012) did to simulate an abiotic planet. In brief, we consider planets with surface pressures of 1 bar, surface temperatures of 288K, and bulk dry atmospheric composition of 10% N 2 , 90% H 2 /CO 2 for the H 2 /CO 2 -dominated cases, respectively. The temperature profile is taken to evolve as a dry adiabat until 160 and 175 K for the H 2 - and CO 2 -dominated cases, respectively, and isothermally thereafter.

The strength of vertical mixing is scaled from that measured in Earth's atmosphere according to the mean molecular mass. The H 2 O concentration at the bottom of the atmosphere is set to 0.01, corresponding to 60% humidity. H 2 , CO 2 , CH 4 , SO 2 , and H 2 S are emitted from the surface at rates corresponding to terrestrial volcanism. The mixing ratio profile of the dominant gases (gases with abundances exceeding 100 ppb) used for the modeling of the H 2 -rich atmosphere on a massive super Earth orbiting an active M dwarf is shown in Fig. 2 (see Appendix A for the mixing ratio profiles used to model the remaining atmospheric scenarios). For further details, including the rationale for these parameters, see Hu et al. (2012).

FIG. 2. Mixing ratio profile of a H 2 -rich atmosphere on an Earth-sized planet orbiting an active M dwarf. Vertical axis represents pressure in units of Pa, and the horizontal axis shows the mixing ratio represented as a percentage of the total atmospheric layer. Figure partially adapted from the works of Hu et al. (2012) and Seager et al. (2013b). Color images are available online.

Stellar irradiation is a key input for photochemical models. We considered instellation corresponding to the Sun (our “Sun-like” case) (Hu et al., 2012) and from the M dwarf GJ1214 (our “active M dwarf” case) (Seager et al., 2013b). The semimajor axes of the planets for the H 2 - and CO 2 -dominated cases are taken to be 1.6 AU and 1.3 AU for the Sun-like case, and 0.042 AU and 0.034 AU for the “active M dwarf” case, corresponding to surface temperatures of 288 K at 0 PH 3 emission.

For a sensitivity test, we also considered a theoretical “quiet” M dwarf simulated by reducing the UV-flux of GJ1214 by three orders of magnitude (corresponding to approximately a factor of 100 less UV radiation than the least active M dwarf known GJ581) (France et al., 2016). See Section 4.1.4 for a discussion of the sensitivity of our results to changes in surface temperature and to low UV irradiation levels.

3.3. Atmospheric spectral simulations

We use the outputs of the photochemical models described above to model observational spectra projected to massive super-Earths with M p = 10 M E and R p = 1.75 R E . We focus on such large massive planets due to the following observational considerations: (1) such planets are easier to detect via radial velocities and transit observations; (2) large planets have larger thermal emission signatures; and (3) massive planets are more likely to retain H 2 -rich atmospheres, which are much easier to characterize in transmission because of their larger scale heights than other atmosphere types.

3.3.1. Transmission and emission spectral calculations

Transmission and thermal emission spectra were simulated with the program SEAS (Simulated Exoplanet Atmosphere Spectra). The projection from Earth-sized to super-Earths is performed by using equivalent techniques to those in the work of Hu et al. (2012); SEAS takes as input a list of molecular mixing ratios as a function of pressure, which, to first order, are invariant to changes in the surface gravity.

The transmission spectrum code calculates the optical depth along the limb path and sums up chords, assuming that the planet atmosphere is homogeneous. The SEAS transmission spectrum code is similar in structure to that described in two studies (Miller-Ricci et al., 2009; Kempton et al., 2017), with the main difference being that SEAS can accept variable mixing ratio inputs, important for super-Earths whose atmospheres are severely impacted by photochemistry. The temperature–pressure profile, including limits and resolution, is specified by the user. Molecular line lists are taken from HITRAN 2016 and ExoMol databases (Gordon et al., 2017; Tennyson et al., 2016), with cross sections calculated with HAPI (Kochanov et al., 2016) and ExoCross (Yurchenko et al., 2018), respectively. The molecular species are chosen by the user, and all molecules in the HITRAN and ExoMol databases are user-selectable options.

The thermal emission code integrates a blackbody exponentially attenuated by the optical depth without scattering (Seager, 2010). The code uses the same input temperature–pressure profile and molecular cross sections as described above.

SEAS considers clouds in the emergent spectra for thermal emission by averaging cloudy and cloud-free spectra (resulting in 50% cloud coverage). We omit clouds or hazes for the transmission spectra model; if the atmosphere is cloudy or hazy at high altitudes, the spectral features in transmission will be muted. Consequently, our calculations represent upper bounds on the magnitude of the transmission features with respect to cloud or haze effects. We discuss the impact of clouds in our modeled transmission spectra in Section 5.2.

The SEAS transmission code has been validated by comparing results with the Atmospheric Chemistry Experiment data set (Bernath et al., 2005) for transmission spectrum and with the MODTRAN spectrum (Berk et al., 1998) for thermal emission spectrum. We also compared results related to this phosphine work with transmission spectra generated by the code described in the work of Hu et al. (2012).

3.3.2. Detectability metric

We study the spectroscopic detectability of phosphine in H 2 - and CO 2 -rich atmospheres, in transmission and emission observation scenarios. In transmission, we compare with the mean transit depth of the planet radius, and in emission with the blackbody curve.

We consider a 6.5 m space telescope, having a quantum efficiency of 25% observing with a 50% photon noise limit. We consider our 1.75 R Earth -planet to be orbiting (1) a 0.26-R sun M dwarf star at 5 pc with an effective temperature of 3000 K; and (2) a Sun-like star. Stellar flux is the source of the noise, combining in-transit and out-of-transit flux noises. The theoretical transmission spectra are based on JWST and its NIRSpec and MIRI instruments (Bagnasco et al., 2007; Wright et al., 2010). To calculate theoretical thermal emission spectra, we consider a secondary eclipse scenario as observed from JWST with the MIRI instruments (both mid- and low-resolution spectrometers). We binned the data to a resolution of R ∼ 10 to increase the significance of detection.

We investigate the detectability of phosphine in exoplanet atmospheres by adapting the detection metrics defined by Seager et al. (2013b) and Tessenyi et al. (2013). The detectability metric is a theoretical metric that uses simulated data.

We first simulate model-independent observational data for all planetary scenarios considered (e.g., using the instrumental constraints of JWST). For an analysis of the transmission spectra models, we then compare the wavelength-dependent transit depth of the planet to the “white-light” transit depth in each wave band (corresponding to the coverage of each instrument).

Phosphine is considered detectable if we can detect opacity at wavelengths corresponding to PH 3 absorption features with statistically significant confidence. To establish the statistical significance of opacity detection in transmission, we assume a simulated spectrum and then assign binned values for the transit depth. We then calculate the wavelength-dependent one-sigma (1-σ) error bar for each binned value (i.e., standard deviation) using only stellar photon noise. The significance of the deviation is calculated with:

where μλ is the wavelength-dependent transit depth of the simulated atmosphere, μλ is the mean transit depth of the white-light averaged waveband, and σ is the uncertainty on the measurement. The uncertainties are estimated based on shot noise. We then assess the detectability of phosphine by simulating a model atmosphere with and without PH 3 and comparing the deviation of each atmosphere from their associated mean. This comparison establishes whether a model atmosphere with PH 3 fits the simulated observational data better than one without PH 3 .

In thermal emission, we use a similar detectability metric to the transmission analysis described above, with the distinction that we calculate the deviation of our modeled atmosphere spectra from its best-fit blackbody continuum (instead of the white light average used for transmission comparisons). The temperature of the blackbody is set by fitting a blackbody curve to the simulated data.

The integration time is a variable parameter in the SEAS models, but features are only considered detectable if they achieve at least a 3-σ interval with 200 observation hours or less (considering 100 h in-transit and 100 h out-of-transit).

3.3.3. Scaling to smaller planets

We performed our spectral simulations for a massive super-Earth planet (M p = 10 M E and R p = 1.75 R E ). In this section, we consider how the prospects for atmospheric characterization scale to smaller, more Earth-sized worlds.

The amplitude of the atmospheric absorption signal in transmission is characterized by Brown (2001):

where R p is the planet radius and M p is the planet mass. This implies that the transmission spectroscopy signal from a 1 R E , 1 M E planet should be twice the signal from the 1.75 R E , 10 M E planet we consider here, and the phosphine surface fluxes required to produce a detectable atmospheric signal should be half of what we model for our super-Earth scenario.§

The amplitude of the thermal emission spectral signal is characterized by:

where λ is the wavelength, B is the blackbody function, and T cont and T atm are, respectively, the brightness temperature in and out of the spectral line under consideration. The above equation implies that the thermal emission signal from a 1 R E , 1 M E planet should be a third of the signal from the 1.75 R E , 10 M E planet we consider here, and the PH 3 surface fluxes required to produce a detectable atmospheric signal should be three times larger than what we model for our super-Earth scenario.

We conclude that spectrally characterizing Earth-sized planets is comparable in difficulty to characterizing super Earth planets, but that characterizing the atmospheres of smaller worlds is somewhat easier in transmission and somewhat harder in emission. These differences do not affect our main conclusions.

4. Results

We find that phosphine as a detectable biosignature gas has three encouraging properties: (1) PH 3 can accumulate to detectable levels in an exoplanet atmosphere, provided it has a high production rate at the planet's surface (Section 4.1); (2) PH 3 has unique spectral features, namely strong bands around 2.7–3.6, 4.0–4.8, and 7.8–11.5 μm, which allow it to be distinguishable from other dominant atmosphere molecules (Section 4.2); and (3) based on the abundances and surface fluxes needed to produce detectable levels of PH 3 , it has no known false positives provided that the planet's surface temperature is below 800 K (Section 4.3). In addition to the above, our results show that, at surface fluxes near the minimum flux necessary to allow for PH 3 detection, a “runaway” effect for PH 3 occurs that drastically alters the atmosphere (Section 5.1).

We present results for the phosphine detection in H 2 - and CO 2 -rich atmospheres, for planets orbiting Sun-like stars and active M dwarf stars.

4.1. Phosphine detection in exoplanet atmospheres

We performed a series of simulations and calculations to explore the prospects for detecting phosphine in an exoplanet atmosphere using transmission and thermal emission spectroscopy. We consider H 2 - and CO 2 -dominated atmospheres, and stellar irradiation environments corresponding to the modern Sun (“Sun-like”) and GJ1214 (“active M dwarf”); see Seager et al. (2013b) for details.

In this section, we first present a set of simulated spectra, both in transmission (Section 4.1.1) and in emission (Section 4.1.2), for atmospheric scenarios with and without phosphine. Here, we predict the minimum abundances required for PH 3 to be detectable in each atmosphere considered. We calculate the PH 3 surface production rates needed for PH 3 to achieve the abundances required for detection in Section 4.1.3, using our photochemical model to simulate the equilibrium distribution of atmospheric gases throughout each atmosphere. Finally, we assess the sensitivity of our results to a variable surface temperature and a host star with low levels of radiation (Section 4.1.4).

4.1.1. Amount of phosphine required for detection via transmission spectroscopy

We find that phosphine is detectable in anoxic atmospheres only if it is able to accumulate to the order of ppb to 100s of ppm, for H 2 - and CO 2 -rich atmospheres, respectively. For comparison, PH 3 is present at the ppt to ppb level on modern Earth. We estimate the photochemical plausibility of PH 3 accumulating to such large abundances in Section 4.1.3.

Unfortunately, even with high concentrations of phosphine in the atmosphere, many tens of hours of JWST time are needed to detect it. The atmosphere mixing ratios, the surface production rates required, and the number of observation hours needed for different planet and host star scenarios are provided in Table 1.

Table 1. Phosphine Mixing Ratios Needed for Detection in Transmission for Different Atmospheric and Stellar Scenarios, with Associated Observation and Surface Flux Requirements Atmospheric scenario Required mixing ratio for detection Minimum observation hours (in-transit + out-of-transit) Associated confidence interval for phosphine detection (σ) H 2 -dominated, Sun-like star 780 ppm 56 3 H 2 -dominated, active M dwarf (Fig. 3) 220 ppb 91 3 H 2 -dominated, active M dwarf 220 ppb 200 4.4 H 2 -dominated, active M dwarf 5 ppb 200 2.5 H 2 -dominated, active M dwarf 0.28% 3 5 CO 2 -dominated, Sun-like star N/A Not detectable N/A CO 2 -dominated, active M dwarf (Fig. 4) 310 ppm 200 2.7 CO 2 -dominated, active M dwarf 7.6% 32 3

Planets with H 2 -rich atmospheres orbiting active M dwarfs require the smallest phosphine abundances for its detection (10s to 100s of ppb; Fig. 3), which can be expected due to their lower mean radical concentrations compared with an oxidized atmosphere (Seager et al., 2013b). H 2 -rich atmospheres also have transmission spectra that are easier to detect than planets with higher mean molecular-weight atmospheres (e.g., CO 2 ) because of their larger scale height, that is, a “puffier” atmosphere.

FIG. 3. Theoretical transmission spectra for a H 2 -rich atmosphere on a 10 M E , 1.75 R E planet with a surface temperature of 288 K orbiting an active M dwarf (1 bar atmosphere composed of 90% H 2 and 10% N 2 ), after 91 h of observation. Top panel: vertical axis shows transit depth of the simulated atmosphere spectra in units of ppm (y axis) and horizontal axis showing wavelength in microns. The orange curve corresponds to the simulated atmosphere spectrum without PH 3 , and the blue curve to an atmosphere spectrum with PH 3 , simulated considering a PH 3 concentration of 220 ppb. Blue error bars correspond to the wavelength-averaged uncertainty within the instrumental waveband; black and gray error bars correspond to the uncertainty of each wavelength bin for atmosphere models with and without PH 3 , respectively. Green and pink shading represent the wavelength coverage of the NIRSpec and MIRI instruments (Bagnasco et al., 2007; Wright et al., 2010). Middle panel: vertical axis shows the statistical significance of detection for two different model atmospheres, with PH 3 (blue) and without PH 3 (orange). Bottom panel: statistical significance of the detection of PH 3 opacities at each wavelength bin. Vertical axis shows size of the statistical deviation between atmosphere models with and without PH 3 (units of σ-interval). In the middle and bottom panels, the horizontal green line represents the 3-σ statistical significance threshold, and the horizontal axes show the individual wavelength bins (microns). For H 2 -dominated atmospheres, the 4.0–4.8 μm spectral band of PH 3 is the most promising feature for detection in transmission. Color images are available online.

For H 2 -rich atmospheres, only the strongest band of phosphine at 4.0–4.8 μm can be detectable (Fig. 3). The other PH 3 features are either too weak or contaminated by other atmospheric molecular species.

For CO 2 -rich atmospheres (Fig. 4), several spectral bands of phosphine show substantial spectral absorptions when compared with atmospheres without PH 3 . Nonetheless, no spectral band of PH 3 can achieve a 3-σ statistical significance even after 200 observation hours when considering a planet orbiting an active M dwarf.

FIG. 4. Theoretical transmission spectra for a CO 2 -rich atmosphere on a 10 M E , 1.75 R E planet with a surface temperature of 288 K orbiting an active M dwarf (1 bar atmosphere composed of 90% CO 2 and 10% N 2 ), after 200 h of observation. Top panel: vertical axis shows transit depth of the simulated atmosphere spectra in units of ppm (y axis), and horizontal axis showing wavelength in microns. The orange curve corresponds to the simulated atmosphere spectrum without PH 3 , and the blue curve to an atmosphere spectrum with PH 3 , simulated considering a PH 3 concentration of 310 ppm. Blue error bars correspond to the wavelength-averaged uncertainty within the instrumental waveband; black and gray error bars correspond to the uncertainty of each wavelength bin for atmosphere models with and without PH 3 , respectively. Green and pink shading represent the wavelength coverage of the NIRSpec and MIRI instruments (Bagnasco et al., 2007; Wright et al., 2010). Middle panel: vertical axis shows the statistical significance of detection for two model atmospheres, with PH 3 (blue) and without PH 3 (orange). Bottom panel: statistical significance of the detection of PH 3 opacities at each wavelength bin. Vertical axis shows size of the statistical deviation between atmosphere models with and without PH 3 (units of σ-interval). In the middle and bottom panels, the horizontal green line represents the 3-σ statistical significance threshold, and the horizontal axes show the individual wavelength bins (microns). In CO 2 -dominated atmospheres, several spectral features of PH 3 have substantial opacities, but no feature achieves a 3-σ statistical significance when compared with the model atmosphere without PH 3 . Color images are available online.

Phosphine is very difficult to detect on planets orbiting Sun-like stars. Planets with CO 2 -dominated atmospheres require longer than 200 observation hours for the detection of PH 3 in transmission spectra, even with the highest surface fluxes considered (3 × 1013 cm−2 s−1). The detection of PH 3 can only achieve a 3-σ statistical significance on planets with H 2 -rich atmospheres for fluxes of 1014 cm−2 s−1 (Table 1); this flux is comparable to the highest recorded PH 3 flux on Earth (above sewage plants) (Devai et al., 1988) and above the values for the biological production of methane, which on Earth corresponds to 1.2 × 1011 cm−2 s−1 (Segura et al., 2005; Guzmán-Marmolejo and Segura, 2015).

The results presented above show that it is possible, but difficult, to detect phosphine in anaerobic atmospheres if it is present as a trace gas. However, if PH 3 production rates increase sufficiently, they outpace the ability of stellar NUV photons to destroy PH 3 , whether via direct photolysis or via generation of radical species. PH 3 may then become a significant component of the atmosphere (e.g., 100s to 1000s of ppm), and its detectability increases dramatically. The PH 3 surface fluxes required to reach this runaway phase (∼1012 cm−2 s−1) are not significantly higher than those required for detection (∼1010–11 cm−2 s−1). For example, with surface production rates only nine times larger than those that produce the atmospheric spectrum shown in Fig. 3, PH 3 reaches the runaway threshold and can be detected with 5-σ statistical significance after only 3 h of observation (Fig. 5). The plausibility of this runaway effect is discussed further in Sections 4.1.3 and 5.1.

FIG. 5. Theoretical transmission spectra for a H 2 -rich atmosphere on a 10 M E , 1.75 R E planet with a surface temperature of 288 K orbiting an active M dwarf (1 bar atmosphere composed of 90% H 2 and 10% N 2 ), at the threshold of the phosphine runaway phase. Top panel: vertical axis shows transit depth of the simulated atmosphere spectra in units of ppm (y axis), after 3 h of observation, and horizontal axis showing wavelength in microns. The orange curve corresponds to an atmosphere spectrum without PH 3 , and the blue curve to an atmosphere spectrum with PH 3 , simulated considering a PH 3 concentration of 0.28%. Blue error bars correspond to the wavelength-averaged uncertainty within the instrumental waveband; black and gray error bars correspond to the uncertainty of each wavelength bin for atmosphere models with and without PH 3 , respectively. Green and pink shading represent the wavelength coverage of the NIRSpec and MIRI instruments (Bagnasco et al., 2007; Wright et al., 2010). Middle panel: vertical axis shows the statistical significance of detection for two model atmospheres with PH 3 (blue) and without PH 3 (orange). Bottom panel: statistical significance of the detection of PH 3 opacities at each wavelength bin. Vertical axis shows size of the statistical deviation between atmosphere models with and without PH 3 (units of σ-interval). In the middle and bottom panels, the horizontal green line represents the 3-σ statistical significance threshold, and the horizontal axes show the individual wavelength bins (microns). Once PH 3 enters the runaway phase, it can be detected after a few hours of observations, through its two strong features in the 2.7–3.6 and 4–4.8 μm regions. Color images are available online.

4.1.2. Amount of phosphine required for detection via emission spectroscopy

We now examine the influence of phosphine in the simulated emission spectra of H 2 - and CO 2 -rich planets orbiting Sun-like stars and active M dwarfs.

Our findings on the amount of phosphine required for detection in thermal emission are similar in transmission, that is, PH 3 can only be detected with many tens of hours of observation time (Table 2). We find that, in emission, PH 3 is detectable in anoxic atmospheres only if it accumulates to at least abundances in the order of ppb; for comparison, PH 3 is present at the ppt to ppb level on modern Earth. The photochemical plausibility of PH 3 accumulating to such large abundances is presented in Section 4.1.3.

Table 2. Phosphine Mixing Ratios Needed for Detection in Emission for Different Atmospheric and Stellar Scenarios, with Associated Observation and Surface Flux Requirements Atmospheric scenario Required mixing ratio for detection Minimum observation hours (in-transit+out-of-transit) Associated confidence interval for phosphine detection (σ) H 2 -dominated, Sun-like star N/A Not detectable N/A H 2 -dominated, active M dwarf (Fig. 6) 220 ppb 131 3 H 2 -dominated, active M dwarf 4 ppm 52 3 CO 2 -dominated, Sun-like star N/A Not detectable N/A CO 2 -dominated, active M dwarf (Fig. 7) 15 ppm 150 3 CO 2 -dominated, active M dwarf 310 ppm 48 3

Detecting an opacity corresponding to an atmospheric spectral feature is much easier than assigning an atmospheric feature to a particular molecule. Our models show that the detection of a CO 2 - or H 2 -rich atmosphere with high statistical significance (>5-σ) is feasible with only a few tens of observation hours. However, the unambiguous attribution of opacity to phosphine requires much longer observation times (Table 2). As an observer, a detection can be considered as an offset to the blackbody curve, but these are only reliable if the blackbody temperature has been accurately estimated. Our detection metric uses a blackbody curve created from a best fit to the simulated observations, which biases toward low temperatures by non-PH 3 absorbers. In reality the temperature of the planet may be obtained through other methods, so our results can be considered a conservative estimate for the minimum PH 3 abundances required for detection.

We find the most detectable spectral region of phosphine in thermal emission is the broad band at 7.8–11.5 μm (Figs. 6 and 7). In emission, planets orbiting an active M dwarf require the smallest PH 3 abundances (100s of ppb to 100s of ppm) to confirm its detection, achieving a 3-σ confidence interval with a minimum of 52 and 48 h of observation, for H 2 - and CO 2 -rich atmospheres, respectively.

FIG. 6. Detectability of phosphine in the emission spectrum of a super Earth exoplanet (10 M E and 1.75 R E ) with a H 2 -rich atmosphere orbiting an active M dwarf, after 131 h of observation. Horizontal axes show wavelength in microns (top) and wave numbers in cm−1 (bottom). Top panel: vertical axes show the flux ratio between the star and the planet; pink and purple lines represent the blackbody curves fitted to the simulated observational data for atmospheres with and without PH 3 , respectively; blue and red curves represent a modeled atmosphere with a PH 3 mixing ratio of 220 ppb and an atmosphere without PH 3 , respectively; black error bars represent the 1-σ uncertainty in the observed data; MRS (yellow shading) and LRS (green shading) represent the coverage of the JWST mid- and low-resolution MIRI instruments, respectively. Middle panel: statistical significance of the detection of an atmosphere with (blue) and without (red) PH 3 when compared with their best-fit blackbody curves, in units of σ-interval; the horizontal green and orange lines represent the 3-σ and 5-σ statistical significance threshold, respectively. Bottom panel: statistical significance of the deviation between an atmospheric model with and without PH 3 ; the horizontal green line represents the 3-σ statistical significance threshold. The detection of PH 3 achieves a 3-σ confidence interval through the high-frequency wing of its strong broad band at 7.8–11.5 μm. JWST, James Webb Space Telescope. Color images are available online.

FIG. 7. Detectability of phosphine in the emission spectrum of a super Earth exoplanet (10 M E and 1.75 R E ) with a CO 2 -rich atmosphere orbiting an active M dwarf, after 48 h of observation. Horizontal axes show wavelength in microns (top) and wave numbers in cm−1 (bottom). Top panel: vertical axes show the flux ratio between the star and the planet; pink and purple lines represent the blackbody curves fitted to the simulated observational data for an atmosphere with and without PH 3 , respectively; blue and red curves represent a modeled atmosphere with a PH 3 mixing ratio of 310 ppm and an atmosphere without PH 3 , respectively; black error bars represent the 1-σ uncertainty in the observed data; MRS (yellow shading) and LRS (green shading) represent the coverage of the JWST mid- and low-resolution MIRI instruments, respectively. Middle panel: statistical significance of the detection of an atmosphere with (blue) and without (red) PH 3 when compared with their best-fit blackbody curves, in units of σ-interval; the horizontal green and orange lines represent the 3-σ and 5-σ statistical significance threshold, respectively. Bottom panel: statistical significance of the deviation between an atmospheric model with and without PH 3 ; the horizontal green line represents the 3-σ statistical significance threshold. The detection of PH 3 achieves a 3-σ confidence interval through its strong broad band at 7.8–11.5 μm. Color images are available online.

Detection of phosphine, in emission, on planets orbiting Sun-like stars is difficult. In these scenarios, the detection of any modeled super-Earth atmosphere cannot achieve a 3-σ confidence interval even with 200 observation hours.

We note that, at sufficiently high phosphine concentrations, our model shows that the wings of the PH 3 absorption features become opaque (e.g., the strong broad band at 7.8–11 μm) and our emission spectra probe the isothermal stratosphere. Consequently, if PH 3 concentrations are high enough, our models show that it is not possible to detect wavelength-dependent opacities due to PH 3 on the basis of emission data alone. At face value, this observation implies a maximum PH 3 concentration and flux, past which it is impossible to detect PH 3 in emission. However, in reality this effect is an artifact of our assumption of an isothermal stratosphere. The stratosphere may have temperature variations, which would facilitate the detection of wavelength-dependent opacity variations due to PH 3 . A coupled climate-photochemistry model that can provide self-consistent temperature–pressure profiles is required to probe this scenario.

4.1.3. Phosphine surface fluxes required for detection

More critical than atmosphere abundances is the surface flux (i.e., the biological production rate) required for phosphine to accumulate to detectable abundance levels in anoxic atmospheres. This quantity plays a key role in determining the efficacy of PH 3 as a biosignature: if the presence of detectable levels of PH 3 in an atmosphere requires surface fluxes of PH 3 that are higher than those which a biosphere could plausible generate, then it is disfavored as a biosignature gas; if, however, PH 3 accumulates to detectable concentrations at fluxes within the range of plausible biological productivity, it is favored as a biosignature gas.

As phosphine moves up the atmosphere, its destruction rate and consequent mixing ratio change, due to the varying levels of radical concentrations and radiation at different altitudes. The dominant PH 3 reaction in H 2 -dominated atmospheres is PH 3 + H. The dominant reaction in CO 2 -dominated atmospheres is PH 3 + O. However, in high-PH 3 atmospheres, H produced from PH 3 photolysis becomes an increasingly important sink for PH 3 , even in CO 2 -dominated atmospheres. PH 3 is unlikely to dissolve into water and condense into aerosols (as ammonia, hydrogen sulfide, and methanethiol are) (Glindemann et al., 2003), meaning rainout is not expected to be a sink.

We use our photochemical model to estimate the minimum surface production flux, P PH3 , for the detectability of phosphine in transmission and emission for a range of planetary scenarios (Table 3). We find that, for planets orbiting active M dwarfs, PH 3 can build to concentrations detectable by transmission and emission spectroscopy if produced at the surface with rates of the order of 1011 cm−2 s−1. We note that PH 3 requires similar surface flux rates in H 2 - and CO 2 -dominant atmospheres to reach detectable abundance levels, even though those correspond to much lower PH 3 concentration requirements in H 2 -rich atmospheres than in CO 2 -rich atmospheres. We speculate that this occurs because UV penetrates deeper into the more transparent H 2 -rich atmosphere, allowing more radical accumulation and more photolysis at depth [see Fig. 4 of Hu et al. (2012)].

Table 3. Phosphine Mixing Ratios Needed for Detection in Transmission and Emission for Different Atmospheric and Stellar Scenarios, As Well As Associated Surface Flux Requirements (P PH 3 [cm−2 s−1]) Atmospheric scenario Required mixing ratio for detection (in transmission and emission) P PH 3 (cm−2 s−1) H 2 -rich planet, Sun-like star 780 ppm (transmission) 1 × 1014 H 2 -rich planet, active M dwarf 5 ppb (transmission) 1 × 1010 H 2 -rich planet, active M dwarf 220 ppb (emission) 1 × 1011 H 2 -rich planet, active M dwarf (PH 3 runaway) 0.28% (transmission) 9 × 1011 CO 2 -rich planet, active M dwarf 310 ppm (transmission) 3 × 1011 CO 2 -rich planet, active M dwarf 15 ppm (emission) 1 × 1011 CO 2 -rich planet, active M dwarf (PH 3 runaway) 7.6% (transmission) 1 × 1012

The phosphine surface fluxes required to generate the detectable levels of PH 3 are large when compared with global PH 3 emissions on Earth but are comparable to the production rates of other major biosignature gases. For comparison, biological CH 4 and isoprene production on Earth are of the order of 1011 cm−2 s−1 (Guenther et al., 2006), where a significant proportion of modern terrestrial CH 4 production is anthropogenic (Houghton, 1995; Segura et al., 2005; Guzmán-Marmolejo and Segura, 2015). As a further comparison, the highest recorded surface flux of PH 3 on Earth is above sewage plants, where PH 3 production reaches 1014 cm−2 s−1 (Devai et al., 1988).

One of our most interesting findings is the existence of a critical phosphine surface production flux, past which PH 3 accumulation is efficient and the atmosphere transitions to a PH 3 -rich state. We term this critical flux the “tipping point.” This effect appears analogous to the “CO runaway” effect identified for early Earth (Kasting et al., 1983, 1984, 2014; Zahnle, 1986; Kasting, 2014). Past the tipping point, PH 3 production outpaces the ability of stellar NUV photons to destroy PH 3 , whether via direct photolysis or via generation of radical species. In this runaway phase, PH 3 can accumulate to percent levels and pervade the atmosphere. In this case, our models show that PH 3 can be detected with observation times reaching under 10 h (Fig. 5). The plausibility of such a PH 3 runaway effect is discussed in Section 5.1.

4.1.4. Sensitivity analysis to temperature and radiation levels

Our approach prescribes a temperature–pressure profile and considers only two possible stellar scenarios (Sun-like stars and active M dwarfs). We conducted sensitivity analyses to assess the dependence of our results on these assumptions.

4.1.4.1. Sensitivity analysis to temperature

In our study, we assumed surface temperatures of 288 K; in reality, worlds with a broad range of temperatures may be habitable. Temperature may affect phosphine concentrations through varying reaction rates of PH 3 with radicals and through changing the concentration of H 2 O in the atmosphere, from which the radical species are largely derived. To test the sensitivity of our results to surface temperature, we calculated PH 3 profiles for CO 2 - and H 2 -rich atmospheres which have detectable concentrations of PH 3 at 288 K, for new surface temperatures of 273 and 303 K. For simplicity, in calculating the dry adiabatic evolution of the lower atmosphere, we approximated the specific heat capacities at constant pressure of CO 2 and H 2 , by their values at 273 K and 303 K (Pierrehumbert, 2010). We adjusted the surface mixing ratio of water vapor to 0.0036 and 0.026, corresponding to the vapor pressures at 273 and 303 K, respectively, for the same 60% humidity assumed at 288 K. Figures 8 and 9 present the results of these sensitivity tests in the case of a H 2 -dominated atmosphere orbiting an M dwarf star.

FIG. 8. Distribution of abundances of atmospheric constituents (top panel) and phosphine (bottom panel) throughout the atmosphere of a H 2 -rich planet orbiting an active M dwarf. Vertical axes represent altitude in units of km, and horizontal axes represent molar concentration. Solid lines and dotted lines correspond to mixing ratios with surface temperatures of 288 and 273 K, respectively. When comparing low temperatures (273 K) to our standard 288 K models, the PH 3 mixing ratio remains mostly unchanged. Color images are available online.

FIG. 9. Distribution of abundances of atmospheric constituents (top panel) and phosphine (bottom panel) throughout the atmosphere of a H 2 -rich planet orbiting an active M dwarf. Vertical axes represent altitude in units of km, and horizontal axes represent molar concentration. Solid lines and dotted lines correspond to mixing ratios with surface temperatures of 288 and 303 K, respectively. When comparing high temperatures (303 K) to our standard 288 K models, the PH 3 mixing ratio remains mostly unchanged. Color images are available online.

We find phosphine abundances to be weakly sensitive to surface temperature. For both CO 2 - and H 2 -dominated atmospheric scenarios, the total PH 3 column varies by a factor of ≤3 relative to the value at 288 K across 273–303 K, with the strongest variation occurring in H 2 -dominated atmospheres. Such variations, while potentially significant for retrievals, do not affect our order-of-magnitude conclusions regarding the detectability of PH 3 .

We attribute this relatively modest variation of phosphine column with temperature to the comparatively small variation of both the PH 3 radical reaction rates and the total water vapor column across this temperature range. Other atmospheric constituents, such as methane, show a much greater sensitivity to lower temperatures than PH 3 . We are unsure why this is the case. One possibility is that PH 3 reaction rates are less sensitive to temperature changes than other atmospheric constituents (e.g., from 288 to 303 K, the rate constant for H + CH 4 increases by a factor of 2.3, whereas the rate constant for H + PH 3 increases by a factor of 1.16). Another possible explanation for CH 4 having a greater sensitivity to temperature than PH 3 is that CH 4 is primarily removed by OH (and therefore most sensitive to H 2 O), whereas PH 3 is primarily removed by O and H (i.e., less sensitive to H 2 O).

We crudely considered the potential impact of high phosphine abundances on the temperature profile of a planet. The greenhouse gas potential of PH 3 is not known (Bera et al., 2009), but it is plausible that a significant accumulation of phosphine on an atmosphere would contribute to an increase in the global temperature since PH 3 is a strong infrared absorber.

To first order, the change in surface temperature due to phosphine can be estimated by calculating the surface temperature required to produce enough outgoing radiation to balance the arriving stellar radiation (Pierrehumbert, 2010). We executed this procedure for an atmosphere with and without PH 3 . We estimate that, if PH 3 accumulates to the abundances required for its detection (see Section 4.1.3), PH 3 can lead to an increase of surface temperature between 10 and 30 K, depending on the atmospheric scenario. Further studies on the greenhouse gas potential of PH 3 are needed to fully explore the impact of its accumulation on the temperature profile of exoplanet atmospheres.

Overall, we conclude that our results are insensitive to variations in surface temperature of ±15 K.

4.1.4.2. Sensitivity to UV irradiation

UV irradiation limits phosphine concentrations through direct photolysis and radical production. We considered the hypothesis that PH 3 would build to higher concentrations on a planet orbiting a star with low UV output, such as a quiet M dwarf, as considered by Domagal-Goldman et al. (2011). To test this hypothesis, we simulated CO 2 - and H 2 -rich planets orbiting a theoretical “quiet” M dwarf.

We constructed our quiet M dwarf model by reducing the instellation at wavelengths <300 nm of our active M dwarf case (corresponding to GJ1214) by a factor of 1000. This corresponds to ∼100 times less UV than GJ581, the quietest M dwarf observed by the MUSCLES survey (France et al., 2016). A truly quiet M dwarf may not exist, as practically all M dwarfs observed to date have at least some chromospheric activity (France et al., 2013, 2016). Our quiet M dwarf case may therefore be considered a theoretical limiting case to study the effect of UV radiation on PH 3 buildup, with the understanding that this limiting case may not exist in reality. We nonetheless note that this limiting case is less extreme than photosphere-only limiting cases considered in past work (Domagal-Goldman et al., 2011; Seager et al., 2013b; Rugheimer et al., 2015).

We find that, for the equivalent surface production rates, phosphine concentrations are two orders of magnitude higher on planets in the quiet M dwarf cases compared with the active M dwarf cases. Low UV emission favors buildup of PH 3 due to lower radical concentrations and photolysis rates. Consequently, as with other proposed biosignature gases, planets orbiting quiet M dwarfs are the best targets for detecting biogenic PH 3 (Segura et al., 2005; Domagal-Goldman et al., 2011; Seager et al., 2013b). We also find that, in planets orbiting a quiet M dwarf, PH 3 is able to enter a runaway phase with two orders of magnitude lower surface fluxes than those required in more active stars (Section 5.1).

Our overall main finding is that, because phosphine is easily destroyed either directly by UV or indirectly by UV-mediated creation of H, O, or OH radicals, a UV-poor environment is favorable for the detection of PH 3 . This result is consistent with past work (Segura et al., 2005; Domagal-Goldman et al., 2011; Seager et al., 2013b). If there are no sufficiently quiet M dwarf stars, we speculate that a UV-poor environment can be created by a UV shield on the planet itself (Wolf and Toon, 2010). Since PH 3 is readily destroyed in an O 2 -rich environment, an ozone UV shield is unsuitable because other oxygen-containing radicals would destroy PH 3 . However, elemental sulfur aerosols generated on planets with high volcanism and reducing atmospheres may provide such a UV shield (Hu et al., 2013). Additionally, if PH 3 fluxes are high enough, they can overwhelm the supply of destructive UV photons and build up to higher concentrations (a runaway effect). For more context and for a comparison with CH 3 Cl, another proposed biosignature gas, see Sections 5.1 and 5.3.

4.2. Phosphine spectral distinguishability

Phosphine's spectral features can be easily distinguished from those of other gases expected to be main components of rocky planet atmospheres. Such gases include water vapor, methane, carbon dioxide, carbon monoxide, ammonia, and hydrogen sulfide (Fig. 10). Ammonia might be present in hydrogen-rich atmospheres (Seager et al., 2013b).

FIG. 10. Comparison of the spectral cross sections of phosphine with other molecular gases at room temperature. Intensity on y axes in a log-scale with units of cm2/molecule, and wavelength represented on the x axes in microns. All cross sections are calculated at zero-pressure (i.e., Doppler-broadened lines only) using the procedure described by Hill et al. (2013). PH 3 , shown in black, is distinguishable from all compared molecules due to its strong bands in the 2.7–3.6, 4–4.8, and 7.8–11 μm regions. Color images are available online.

The infrared spectrum of phosphine has three major features: 2.7–3.6, 4–4.8, and 7.8–11 μm, corresponding to polyad (P) numbers 3, 2, and 1, respectively (Sousa-Silva et al., 2013). The 2.7–3.6 μm region (P = 3) is dominated by a hot band and an overtone band, both associated with the symmetric bending mode of PH 3 , and six additional combination bands. The 4–4.8 μm region (P = 2) is dominated by both the fundamental symmetric and asymmetric stretching bands of the PH 3 molecule. The P = 2 feature is also where several weaker, bending overtone and combination bands occur, combining with the fundamental bands to result in the strongest overall spectral feature for PH 3 . When compared with water, methane, ammonia, and hydrogen sulfide (but not CO 2 ), the P = 2 feature is uniquely attributable to PH 3 (Fig. 10). In the 7.8–11 μm region (P = 1), the fundamental symmetric and asymmetric bending modes, as well as hot bands, combine to produce a strong and broad absorption feature. The P = 1 PH 3 feature overlaps with the ammonia spectrum but is easily distinguishable from the remaining molecules in this comparison (Fig. 10).

For more detailed comparisons focusing on the 2.7–5 and 7.8–11.5 μm regions, see Appendix B.

The strongest band of phosphine, in the 4–4.8 μm region, is particularly salient when comparing PH 3 with all available spectra of volatile molecules (Fig. 11). However, the second strongest feature of PH 3 , the broad band in the 7.8–11 μm region, is easily obscured by other gases as it absorbs in a heavily populated wavelength region, where many molecules have strong fundamental rovibrational modes.

FIG. 11. Comparison of the spectral cross sections of phosphine (orange) with all the available cross sections for molecules that are volatile at room temperature (Lemmon et al., 2010). Intensity on y axes in a linear scale representing absorbance (normalized to 1), and wavelength represented on the x axes in microns, with the spectral range constrained to 2.5–18.5 μm for fair comparison (many molecules have incomplete spectra beyond this region). Opacity for all molecules is plotted at 1% so that heavily populated regions are highlighted. All cross sections are calculated with SEAS, using molecular inputs from NIST (Linstrom and Mallard, 2001) and ExoMol (Tennyson et al., 2016). The strongest band of PH 3 (4.0–4.8 μm) is easily distinguishable from all other gases, but the broad band at 10 μm can become obscured by other molecules. SEAS, Simulated Exoplanet Atmosphere Spectra. Color images are available online.

It is worth noting that, of the 534 molecules for which there are available spectra, only a few dozen have been adequately measured or calculated, and consequently their spectra should be considered preliminary. Furthermore, there are thousands of volatile molecules that could contribute to an atmospheric spectra (Seager et al., 2016) for which there is no available spectral data, so further studies are required to reveal the full extent of the spectral comparison highlighted in Fig. 11 (Sousa-Silva et al., 2018, 2019).

4.3. Phosphine false positives

On Earth, the only significant amounts of phosphine found in the atmosphere are produced anthropogenically or biologically (Sections 2.1 and 2.2). The formation of PH 3 on temperate rocky planets is thermodynamically disfavored, even in high-reducing environments, unlike the abiological production of methane or hydrogen sulfide.

In thermodynamic equilibrium, phosphorus can be conservatively expected to be found in the form of phosphine only at T > 800K, and at P > 0.1 bar (Visscher et al., 2006), and only in environments rich in H 2 , which is why PH 3 has been detected in Jupiter and Saturn, where these extreme temperatures occur (in the deep layers of the atmosphere). In H 2 -poor environments, much higher temperatures and pressures are required for PH 3 to be thermodynamically favored. We also note that the critical temperature of water is 647 K so there are no surface conditions that favor both PH 3 production and allow for the presence of liquid water. Consequently, in a temperate rocky planet, it is implausible that PH 3 can be produced without biological intervention, so its detection in such an environment is a promising indication of biological activity. We summarize below the potential false-positive scenarios for PH 3 as a biosignature gas and their expected impact on the global concentrations of PH 3 .

Overall, nonbiological phosphine formation is not favored on temperate rocky worlds, and no abiotic pathways can produce PH 3 with the production rates necessary for its detection on habitable exoplanets. We therefore conclude that, in contrast to molecules such as ammonia and methane, a detection of PH 3 on a temperate exoplanet is likely to only be explained by the presence of life.

4.3.1. Formation from phosphite and phosphate

We considered the hypothesis that phosphine could be formed geochemically as a “false positive” by reduction of phosphate or phosphite to PH 3 . Phosphate is a dominant form of phosphorus on Earth. Phosphite is much less abundant but was detected in ground water and in mineral deposits (Han et al., 2012, 2013; Yu et al., 2015) where it is likely to be the result of biological activity (Bains et al., 2019a). Phosphite is also postulated to have been much more abundant on early anoxic Earth (Pasek, 2008; Pasek et al., 2013; Herschy et al., 2018). We calculated the Gibbs free energy of formation of PH 3 from both phosphate and phosphite under geochemical source conditions at neutral pH, for T = 273 and 413 K and pH 2 = 10−6 and 1 bar. In all cases, the formation of PH 3 was thermodynamically disfavored; see Appendices C and D for details. We conclude that PH 3 formation from phosphate or phosphite is unlikely in the absence of a biological catalyst; for more details on the thermodynamic plausibility of the reduction of phosphites or phosphate into PH 3 , see Bains et al. (2019a).

Phosphite can disproportionate to phosphine at T ≳ 320K and acidic pH (Bains et al., 2019a), raising the possibility that “black smoker” hydrothermal systems (T ≤ 678K, pH = 2–3) (Martin et al., 2008) might generate phosphine (Appendix D). Such systems do not dominate volcanic emission on Earth, leading us to propose they would be a negligible contributor on Earth-analog worlds. However, if a world had global, hot, acidic oceans (e.g., due to very high pCO 2 ), then the theoretical possibility of abiotic phosphine production exists, although likely only in the presence of high H 2 concentrations, very low pH and within a very hot temperature band (Appendix D) (Bains et al., 2019a). Given that these oceans would be unlikely to have pH values below 4 (carbonic acid has a pH of 3.6) and PH 3 formation is only favored at pHs closer to 2, we consider this scenario possible but implausible.

4.3.2. Lightning

We also considered the possibility of phosphine production by lightning. Lightning discharges even in highly reducing atmospheres produce only negligible amounts of reduced phosphorus species, including PH 3 , and are very unlikely to provide high flux sources of PH 3 globally. A few studies have examined the production of reduced phosphorus species from phosphate as a result of simulated lightning discharges in laboratory conditions (Glindemann et al., 1999., 2004); only a very small fraction of the phosphorus was reduced to PH 3 through this process, even in highly reduced atmospheric conditions. Similarly, a mineral fulgurite—a glass resulting from lightning strikes—was also proposed as a potential source of PH 3 given that it could, in principle, contain reduced phosphorus species (Pasek and Block, 2009). However, these sources are rare and localized; they would have minimal impact on a global scale. We are not aware of kinetically favored reactions that would promote the conversion of the thermodynamically favored phosphate to PH 3 .

4.3.3. Volcanism

Phosphine is not known to be produced by volcanoes on Earth.

Calculations on the production of phosphine through volcanism on a simulated anoxic early Earth showed that only trace amounts of PH 3 can be created through this avenue; the predicted maximum production rate is 102 tons per year (Holland, 1984), which corresponds to ∼104 cm−2 s−1. We note that the estimation of the maximum production of PH 3 through the volcanic processes reported by Holland (1984) is made under the assumption of a highly reduced planet, which provides favorable conditions for PH 3 volcanic production. The volcanic production of PH 3 in other planetary scenarios is even more unlikely. We estimate that the maximum production of PH 3 by volcanoes in any planetary scenario, even H 2 -rich atmospheres, is at least seven orders of magnitude lower than the surface fluxes required for detection (see Section 4.1.3).

4.3.4. Exogenous delivery

Finally, we considered the possibility of exogenous meteoritic delivery as a source of reduced phosphorus species that could lead to the abiotic production of phosphine. Reduced phosphorus species can be found in the meteoritic mineral schreibersite (Pech et al., 2011). Schreibersite is (Fe,Ni) 3 P, which is present in iron/nickel meteorites (Geist et al., 2005); it is not present in stony or carbonaceous bodies. The current accretion rate of meteoritic material to Earth is of the order of 20–70 kilotons per year (Peucker-Ehrenbrink, 1996). We calculated the maximum PH 3 production from these sources as follows: considering that ∼6% of the meteoritic material is iron/nickel (Emiliani, 1992) and such meteorites contain an average of 0.25% phosphorus by weight (Geist et al., 2005), and working under the conservative assumption that the totality of the phosphorus content could be hydrolyzed to PH 3 , these meteors would deliver a maximum of ∼10 tons of PH 3 to Earth every year.** Therefore, the contribution from meteoritic sources to the global average PH 3 production rates is still negligible. The above calculations are also in agreement with previous estimations of the phosphine production through meteoritic delivery, which were also found to be negligible (Holland, 1984).

5. Discussion

We find that phosphine is a promising marker for life if detected on a rocky exoplanet. On Earth, PH 3 is naturally associated exclusively with anaerobic life and is not expected to have any significant false positives for life on temperate exoplanets. Our models find that, if produced at sufficiently high surface fluxes, PH 3 can accumulate in planetary atmospheres to detectable abundances. Here, we discuss the photochemical impact of high abundances of PH 3 in an exoplanet atmosphere (Section 5.1). We then describe the known limitations of our calculations (Section 5.2) and expand on alternative methods of detecting PH 3 in exoplanet atmospheres (Section 5.3). We summarize our findings in section 5.4.

5.1. Phosphine “tipping point” and its impact on the atmosphere

Our models show that, with global phosphine surface fluxes comparable to those found locally in anoxic ecosystems on Earth, PH 3 can have a significant impact on planets with anoxic atmospheres. We calculate that PH 3 becomes detectable on anoxic planets where it is emitted at the surface with fluxes greater than 1011 cm−2 s−1. Consequently, PH 3 is accessible to remote detection only if it is a substantial product of the biosphere, emitted in quantities similar to CH 4 and isoprene on Earth (Guenther et al., 2006). We note that the maximum recorded surface flux of PH 3 on Earth is 1014 cm−2 s−1 (Devai et al., 1988). PH 3 may be emitted by life at these detectable surface fluxes as a product of primary metabolism, such as CH 4 , on warm acidic worlds, or as a secondary metabolite, such as isoprene (Bains et al., 2019a, 2019b).

A surprising result of our models was that, once the phosphine surface flux reaches a tipping point (e.g., >9 × 1011 cm−2 s−1 for planets orbiting active M dwarfs), PH 3 enters a runaway phase and begins to drastically change the atmosphere (Fig. 5). This phase appears analogous to the CO runaway described for early Earth and reviewed in the work of Kasting et al. (2014). In this runaway phase, PH 3 production outpaces the ability of stellar NUV photons to destroy PH 3 (whether via direct photolysis or via generation of radical species), and modest increases in PH 3 flux lead to dramatic increases in PH 3 accumulation in an atmosphere.

With fluxes of 1012 cm−2 s−1 (10 times higher than the minimum required for detection), phosphine can approach percent concentrations and would be detectable with just less than 10 h of observation (Table 1 and Fig. 12). In this runaway phase, PH 3 can affect the concentrations of other atmospheric constituents; for example, in CO 2 -dominated atmospheres, the H 2 concentration increases dramatically in PH 3 runaway, presumably from H 2 generated by PH 3 destruction. This raises the possibility that enhanced H 2 concentrations may be used to confirm PH 3 detections. In summary, PH 3 is readily detectable if it is produced at rates about an order of magnitude higher than methane or isoprene on Earth.

FIG. 12. Distribution of abundances of major atmospheric constituents (top panel) and phosphine (bottom panel) throughout the atmosphere of a H 2 -rich planet orbiting an active M dwarf. Solid and dashed lines show molecular abundances immediately below and above the PH 3 runaway phase, respectively. The x axis shows abundance concentrations, and the y axis shows altitude in units of km. The scavenging effect of PH 3 leads to a decrease in O and OH and, to some extent, H radicals in the atmosphere. Consequently, both PH 3 and other trace gases (e.g., H 2 S) are able to accumulate to larger abundances once PH 3 enters the runaway phase. Color images are available online.

Phosphine can affect the spectrum of a rocky planet atmosphere, even at concentrations somewhat below detectability, by driving down radical concentrations due to its intense reactivity with these molecules, effectively becoming a scavenger in the atmosphere. This affects the concentrations of other atmospheric constituents, such as methane. Domagal-Goldman et al. (2011) described a comparable effect in models of organosulfur volatiles, which detectably altered the CH 4 and C 2 H 6 abundances despite being themselves undetectable. Domagal-Goldman et al. (2011) reported that elevated C 2 H 6 /CH 4 ratios could be diagnostic of high organosulfur flux, and hence a biosignature. Similarly, it may be possible to use the indirect effects of high PH 3 flux to infer the presence of PH 3 even if it is not directly detectable; more detailed measurements of the reaction kinematics of PH 3 and its by-products, as well as more sophisticated atmospheric modeling are required to explore this possibility.

5.2. Model limitations and assumptions

We have assessed the potential of phosphine as a biosignature gas using a set of sophisticated photochemical and radiative transfer models. Nonetheless, given the complexity of simulating both atmospheric composition and subsequent spectral observations, many approximations and assumptions were made. Below is a brief discussion of the major limitations of the work performed here.

5.2.1. Clouds

We assumed cloudless skies for the simulation of exoplanet transmission spectra. To estimate how clouds might affect our results, we re-ran our models for phosphine abundances in the detectable range considering cloud decks at various altitudes, with coverage ranging from 10% to 100%. As expected, the detectability of PH 3 is reduced with the introduction of cloud coverage; for example, for planets with a H 2 -dominated atmosphere orbiting an active M dwarf, and with PH 3 surface fluxes of 1011 cm−2 s−1, the detectability of PH 3 reduced by up to a factor of 10 (from a cloudless model to a model with full cloud coverage at 100 Pa). However, once PH 3 reaches the tipping point (see Section 5.1), it becomes sufficiently abundant in the upper troposphere to be mostly unaffected by the presence of clouds.

5.2.2. Reaction networks and haze formation

In this work, we have focused on the reactivity of phosphine with the dominant radical species O, H, and OH. The photochemistry of PH 3 with radicals originating from other more exotic atmospheric species, although likely to be rare, is insufficiently studied. For example, photochemistry of PH 3 and hydrocarbons, through UV radiation, could lead to the formation of complicated alkyl-phosphines (Guillemin et al., 1995, 1997), and in consequence increase the probability of hazes. Inclusion of these reactions would increase PH 3 destruction rates and hence increase the required surface fluxes for detection; however, since these species are not expected to be dominant radicals, the effect of their inclusion would be minor and should not affect our results. In contrast to the formation of hydrocarbon and sulfur-based hazes that have previously been thoroughly addressed (Domagal-Goldman et al., 2011; Arney et al., 2017), there is very little work on phosphorus-based hazes. The formation of such hazes is possible in theory (Guillemin et al., 1995, 1997; Pasek et al., 2011), and early laboratory experiments implied the possibility of formation of such organophosphine hazes in planetary atmospheres, but further studies are needed to properly address the plausibility and impact of organophosphine haze formation and its associated potential as a PH 3 sink. Our photochemical model will continue to update whenever we are able to expand our reaction networks.

We have neglected the formation of organic hazes in this work. There is evidence that such hydrocarbon hazes occurred on Earth, due to transient high levels of methane (Izon et al., 2017; Zahnle et al., 2019). Organic hazes are predicted to form at [CH 4 ]/[CO 2 ] > 0.12 for M dwarf stars, and [CH 4 ]/[CO 2 ] ≥ 0.2 for Sun-like stars (Arney et al., 2016, 2017). A methanogenic biosphere producing high CH 4 fluxes is required to generate such high ratios; in our work, we have not considered such a biosphere. The net effect of hazes would be to facilitate phosphine buildup through attenuation of photolytic UV. However, these same hydrocarbon hazes would also cloak some of the PH 3 spectral features, although primarily not in the wavelength bands where PH 3 is a strong absorber.

We did not include in our models the recombination of phosphine via PH 2 + H → PH 3 . The rate constant for this reaction is 1.1 × 10−10 cm3 s−1 at 288 K (Kaye and Strobel, 1984). The rate constant for H attack on PH 3 is 3.3 × 10−12 cm3 s−1 at 288 K (Arthur and Cooper, 1997). Consequently, if [PH 2 ]/[PH 3 ] ≥ (3.3 × 10−12 cm3 s−1)/(1.1 × 10−10 cm3 s−1) = 0.03, then reactions with H can reform PH 3 as fast as it is destroyed by H-attack, and substantially lower the PH 3 surface fluxes that are required for PH 3 to accumulate in the atmosphere. Detailed photochemical modeling is required to constrain whether such high [PH 2 ]/[PH 3 ] is plausible, or whether other sinks will suppress [PH 2 ]. We note that, in models of Saturn's atmosphere, [PH 2 ]/[PH 3 ] << 0.03 (Kaye and Strobel, 1984); if the atmospheres of terrestrial H 2 -dominated exoplanets behave similarly, this recombination mechanism will not be able to significantly replenish PH 3 .

Finally, we note that our results reflect the prediction of atmospheric models that anoxic atmospheres should have much higher mean radical concentrations than modern Earth (Hu et al., 2012). Due to their reactivity, one might expect concentrations of radicals to be suppressed even in anoxic atmospheres, as OH is on Earth. While our reaction networks include all known relevant atmospheric chemistry, it is possible that there are chemical reactions, which are relevant to anoxic temperate terrestrial planets, but which have not been considered in the context of the Solar System and hence are not included in the reaction compendia we use (e.g., the NIST database (Linstrom and Mallard, 2001) and the JPL compendium (Sander et al., 2011).

Due to the incompleteness of our reaction network, it is likely that, when phosphine is destroyed, additional radicals are created that are not considered in our models. While PH 3 is a trace gas, this omission should have negligible consequences, as the calculation of the infinite series due to radicals tends to converge quickly. However, when PH 3 approaches the runaway phase and becomes the dominant radical sink the atmosphere, our scenarios become a low-radical regimen. In reality, the radical production would not stop, and the transition to a runaway scenario may be slower than we predict. If radical concentrations were overestimated in our model, then PH 3 can build to detectable levels with lower surface fluxes. If, at high concentrations of PH 3 , the intermediate radical production from its destruction has been underestimated, then PH 3 can only enter a runaway phase with higher surface fluxes than those calculated here. Detailed studies of the chemical reaction networks of anoxic planets are required to explore these possibilities.

5.2.3. Prescribed temperature–pressure profiles

Our photochemical model uses prescribed temperature–pressure profiles (Appendix A) that are isothermal above the stratosphere. In reality, this is an oversimplification; one of the consequences of assuming that there is no temperature inversion at high altitudes is the underestimation of the detectability of phosphine in our simulated emission spectra (Section 4.1.2). We also did not couple the potential heating effect from PH 3 dissociation in our photochemical model, so we performed a sensitivity analysis to small changes in temperature (±15 K) and found that our main conclusions remain unchanged (Section 4.1.4). However, if PH 3 fluxes exceed the “tipping point” and PH 3 enters a runaway phase, it may be possible for temperatures to increase beyond the maximum 303 K we consider here. In the most extreme scenario, a PH 3 runaway might trigger a runaway greenhouse state, potentially rendering the planet uninhabitable. A coupled climate-photochemistry model is required to thoroughly investigate this scenario.

5.2.4. Phosphine sinks

Our photochemistry model assumes a dry deposition velocity of zero for phosphine, that is, no consumption of PH 3 by surficial geochemistry or biology. Apart from its efficient oxidation by atmospheric components, there are no other known significant PH 3 sinks on Earth. It is possible, however, that other planets may have PH 3 deposition pathways that we cannot account for. For example, given the opportunistic nature of biology, and the fact that, at least on Earth, phosphorus is a growth-limiting nutrient, life might use any excess atmospheric PH 3 as source of phosphorus. It is plausible that anaerobic life on other planets will not only just produce PH 3 but also reabsorb it from the atmosphere; in these scenarios, biology would slow down the PH 3 accumulation in the atmosphere leading to a dampening, and possible avoidance, of a runaway PH 3 effect (Section 5.1). In extreme cases, for example, if biological production of PH 3 can equal its reabsorption, life's recycling of biogenic PH 3 might entirely prevent its accumulation in the atmosphere of an exoplanet.

To estimate the effect of a potential sink for phosphine, we tested the variability of our results to a non-zero deposition velocity for PH 3 . Deposition velocity tests show that, for a comparable deposition rate to CO and O 2 (10−4 cm−1) (Harman et al., 2015), concentrations of PH 3 varied by a factor of <2, which is not enough to affect detectability. We note that this is a conservative estimate given that the only plausible PH 3 sinks are biological.

5.2.5. Phosphorus availability

We considered whether the availability of phosphorus in the crust of a planet could be a limiting factor for the accumulation of phosphine in the atmosphere. An estimate of the total phosphorus within Earth's crust shows that, if all the phosphorus was to be converted to PH 3 , it would produce approximately twice as many PH 3 molecules as the total number of all molecules belonging to all gases present in the atmosphere of the modern Earth (Yaroshevsky, 2006). We conclude that, in principle, the total mass of phosphorus in a planetary crust does not limit the development of a high-PH 3 atmosphere.

5.3. Alternative detection methods beyond JWST

For the phosphine detectability calculations in this work, we have considered observations from a JWST-like telescope, with a 6.5 m diameter telescope mirror operating within 50% of the shot noise limit and a quantum efficiency between 20% and 25%. The integration time is assumed to be under 200 h for all atmospheric scenarios. For comparison, the cryogenic lifetime of JWST is 5 years, which is equivalent to an integration time of 100 h for a planet orbiting the habitable zone of an M dwarf star. The spectral resolution of JWST is R = 100 at 1–5 μm and R = 160 at 5–12 μm, which is more than necessary for distinguishing between PH 3 and other dominant gases in the atmosphere (see Section 4.2).

We considered the possibility of detecting phosphine in anoxic atmospheres using alternative telescopes to JWST. Missions such as TPF-I (Lawson et al., 2008), Darwin (Fridlund, 2000), Origins Space Telescope (OST) (Battersby et al., 2018), Habitable Exoplanet Imaging Mission (HabEx) (Gaudi et al., 2018), Large UV Optical Infrared Surveyor (LUVOIR) (Roberge, 2019), and the 30-m class of ground telescopes (Johns et al., 2012; Tamai and Spyromilio, 2014; Skidmore et al., 2015) could also be able to characterize atmospheres of temperate planets in wavelength regions where PH 3 is spectrally active.

TPF-I was intended to be a nulling interferometer with four 4 m diameter telescopes flying in formation with a baseline range of 40–100 m and operating at 6.5–18 μm with a spectral resolution of 25–50. Darwin was planned as set of 3–4 m diameter telescopes flying in a nulling interferometer configuration. For both instruments, only the nearest (∼4 pc) M dwarf star habitable zones would be accessible. Winters et al. (2019) estimated that there are only 22 M dwarf stars candidates that could have planets within their habitable zone suitable for atmospheric characterization. TPF-I and Darwin are currently canceled, but similar telescopes [e.g., the proposed Large Interferometer For Exoplanets, or LIFE (Quanz et al., 2019)] may one day be commissioned that perform in a similar interferometer formation and could have the capability to identify phosphine on temperate exoplanets orbiting M dwarfs.

The OST (Battersby et al., 2018) has a large wavelength coverage (2.8–20 μm) and will be able to provide atmospheric spectra for planets orbiting K and M dwarf stars, through transmission and secondary eclipse observations. The OST will have the sensitivity and coverage to detect many spectral signatures of potential biosignature gases, including phosphine.

The HabEx (Gaudi et al., 2018) and the LUVOIR (Roberge, 2019) are planned as powerful telescopes that could launch in the coming decades and are focused on the detection of potential biosignature gases. Both cover a wavelength region where phosphine is spectrally active (UV to near-infrared), although not a particularly strong absorber (1.28–1.79 μm) (Sousa-Silva et al., 2014). We have not modeled atmospheres with PH 3 in an HabEx/LUVOIR scenario, but these telescopes' high-contrast spectroscopy may allow f