Author's personal copy
Icarus 212 (2011) 751–761
Contents lists available at ScienceDirect
Icarus
journal homepage: www.elsevier.com/locate/icarus
Phosphorus chemistry on Titan
Matthew A. Pasek a,⇑, Olivier Mousis b, Jonathan I. Lunine c,d
a
Department of Geology, 4202 E Fowler Ave, SCA 528, Tampa, FL 33620, United States
Université de Franche – Comté, Observatoire de Besançon, Institut UTINAM, CNRS/INSU, UMR 6213, France
Department of Planetary Sciences, University of Arizona, 1629 E University Blvd, Tucson, AZ 85721, United States
d
Dipartimento de Fisica, Università di Roma ‘‘Tor Vergata’’, Via della Ricerca Scientifica, 1, 00133 Rome, Italy
b
c
a r t i c l e
i n f o
Article history:
Received 2 August 2010
Revised 1 January 2011
Accepted 19 January 2011
Available online 1 February 2011
Keywords:
Astrobiology
Accretion
Prebiotic environments
a b s t r a c t
Phosphorus is a key element in biology and acts in many critical biochemical functions. The chemistry of
phosphorus in the outer Solar System has not yet been quantified, hence the astrobiological relevance of
phosphorus to possible life on places like Titan is unconstrained. We evaluate phosphorus chemistry on
Titan using a combination of modeling and laboratory techniques. We show that phosphorus chemistry
on Titan consists of exogenous phosphates and reduced oxidation state phosphorus compounds, and
accretionary phosphine. Accretionary phosphorus is shown to be delivered primarily by rocks and ices
in the saturnian sub-nebula, and heating during accretion concentrates phosphine in the crust of Titan.
The exogenous compounds are capable of performing biologically-relevant chemistry, however they
are active only in environments with substantial liquid water, either pure, or as a mixture with NH3 or
nitrile compounds. In contrast, we show that phosphine is soluble in methane and ethane on Titan’s surface, hence phosphine likely participates in the hydrocarbon cycle on Titan. The lack of mobility of phosphate compounds on Titan’s surface suggests that if life is present on Titan, it must have a fundamentally
different biochemistry than does terrestrial life.
! 2011 Elsevier Inc. All rights reserved.
1. Introduction
The elements carbon, nitrogen, oxygen, hydrogen, and phosphorus comprise the biogenic elements and are essential to life as we
know it. Of these, the chemistry of phosphorus (P) is perhaps the
least well understood in the satellites of the outer Solar System,
largely due to its scarcity and the lack of a major phase sufficiently
volatile to appear on the surfaces of the icy moons or the atmosphere of Titan. Phosphorus plays major roles in the form of phosphate biomolecules in replication and information as RNA and
DNA, metabolism as ATP, NADPH, and other coenzymes, and structure as phospholipids. Hence understanding its chemistry in the
outer Solar System is essential to gauging various environments
there as likely to host life, either terrestrial-type or exotic life
(COEL, 2007).
The formation of critical phosphate biomolecules typically requires energetic conditions or activated P compounds (e.g., Keefe
and Miller, 1995; Pasek, 2008). Recent work has shown reduced
oxidation state P compounds may fill the role of activated P compounds (e.g., Pasek and Lauretta, 2005; Pasek et al., 2007; Bryant
et al., 2010). Other P compounds capable of forming organic phosphate compounds include an assortment of phosphate minerals
⇑ Corresponding author.
E-mail addresses: [email protected] (M.A. Pasek), [email protected]
(O. Mousis), [email protected], [email protected] (J.I. Lunine).
0019-1035/$ - see front matter ! 2011 Elsevier Inc. All rights reserved.
doi:10.1016/j.icarus.2011.01.026
(Saladino et al., 2006; Costanzo et al., 2007) heated in the presence
of condensing agents.
While P chemistry in rocky planets and asteroids is relatively
well constrained, it is not a good guide to the details of such chemistry in the outer Solar System. The chemistry of the outer Solar
System and the inner Solar System vary substantially as a result
of differences in oxygen fugacity, due in part to the presence of a
water condensation front, or snow line in the 2–5 AU region during
the existence of the gaseous disk (Lunine et al., 2000; Pasek et al.,
2005). Nonetheless, consideration of P chemistry is crucial for
gauging the potential for life of promising astrobiological targets
in the outer Solar System such as Saturn’s moon Titan (Shapiro
and Schulze-Makuch, 2009).
The prebiotic and biotic potential of Titan have been discussed
variously with respect to the formation of nucleobases (Pilling
et al., 2009), of amino acids (e.g., Neish et al., 2010), the occurrence
of acetylene-based metabolisms (McKay and Smith, 2005; Abbas
and Schulze-Makuch, 2002; Schulze-Makuch and Grinspoon,
2005), and of methane-based life in general (Benner et al., 2004).
The formation of life as we know it requires the presence of either
phosphate or energetic P compounds. Could these compounds in
fact form on Titan? Would life on Titan be dependent on phosphate
compounds, like it is here on the Earth? Or are the P compounds
present on Titan’s surface so different as to preclude the occurrence of a terrestrial-type form of life in appropriate Titan environments such as a subsurface liquid water ocean? And what might
Author's personal copy
752
M.A. Pasek et al. / Icarus 212 (2011) 751–761
the role of P be in an exotic form of life dependent on a liquid medium other than water?
To this end, we set out to determine the chemistry of P on Saturn’s moon Titan. We consider the primary sources of P on Titan,
how these primary forms of P are modified by planetary processes,
and how P compounds are transported on or to Titan’s surface. We
use experimental analyses to expand on P transport on Titan.
While focusing on Titan, many of the processes described here
are also applicable to planetary bodies elsewhere.
2. Types of phosphorus compounds
Phosphorus has one of the large ranges of oxidation states of
any element, from +5 in phosphate to !3 in phosphine (PH3). We
consider sources of five different varieties of P that have been observed in either terrestrial or cosmochemical samples, and discuss
how these compounds may evolve when placed in a variety of conditions (Fig. 1).
Phosphates consist of PO4 tetrahedral in which the oxidation
state of P is nominally +5. Phosphates are the dominant P carrier
in the Earth’s crust and in carbonaceous meteorites. Phosphates
minerals include apatite with formula Ca5(PO4)3(OH,F,Cl), and struvite with formula MgNH4PO4"6H2O. Included with phosphates are
condensed phosphates which are polymers of phosphate, with the
dimer, pyrophosphate, P2 O4!
7 , as the simplest form.
On the Earth, phosphates typically occur with felsic silicates,
and in the Solar System phosphates are formed via the oxidation
of phosphides in meteorites (Olsen and Fredriksson, 1966; Pasek,
2008). Condensed phosphates are typically formed through biologic processes (Kornberg, 1950), although a host of abiotic condensation (i.e., loss of water) processes can also lead to their
formation (see review by Keefe and Miller, 1995). The oxidation
of phosphides can also form pyrophosphate and larger phosphates
through a radical reaction pathway (Pasek et al., 2007). Phosphates
are typically poorly soluble in water, and although there do exist
some soluble sodium and potassium phosphates, in general most
phosphates do not readily dissolve in sea water. The presence of
chelating agents such as acetate, CH3COO! (Schwartz, 1971), or
bicarbonate, HCO!
3 (Kakegawa et al., 2002), can liberate phosphate
to abundances of up to 10!2 moles per liter (M) in solution.
Reduced P oxides include species in which the oxidation of P is
less than the +5 found in phosphates, and include phosphite
!
(HPO2!
3 , +3 oxidation state), hypophosphite (H2 PO2 , +1 oxidation
state), and hypophosphate (P2 O4!
,
+4
oxidation
state).
These forms
6
of P are exceptionally rare in the terrestrial system, and have only
been reported to occur twice in geological samples (Pasek and
Block, 2009; Pech et al., 2009). These compounds are formed by
the reactions of phosphides such as schreibersite with water
(Pasek and Lauretta, 2005; Bryant and Kee, 2006; Pasek et al.,
2007; Bryant et al., 2009).
Reduced P oxide compounds are significantly more soluble that
phosphates, and have greater reactivity. In warm (by terrestrial
standards) environments they will react to form phosphonate
compounds (e.g., Bryant et al., 2010). In oxidizing environments
phosphite and hypophosphite will oxidize to form phosphates
including condensed phosphates (Pasek et al., 2008).
Phosphonates are compounds of the formula R ! PO2!
3 , where R
is an organic radical. The oxidation state of P in these compounds is
nominally +3, as they are organic analogs of phosphite. A number
of phosphonates are found on the Earth, and are formed by
biological processes (Dyhrman et al., 2009), and include aminoethylphosphonate and acetylphosphonate. They are known to occur in the carbonaceous meteorite Murchison, as methyl-, ethyl-,
and propyl-phosphonates (Cooper et al., 1992). These meteoritic
phosphonates may have been formed by corrosion of phosphide
minerals (de Graaf et al., 1995), or by gas-phase reactions in the
ISM (Gorrell et al., 2006). Like reduced P oxides, phosphonates
are more soluble than phosphates in water, and presumably would
be so even in non-polar organic solutions.
Phosphides are metal-phosphorus minerals lacking oxygen,
hence their oxidation state is typically stated as ranging from 0
to !3. Schreibersite, (Fe,Ni)3P, is the major phosphide mineral that
is found in meteorites, in comets and interplanetary dust (e.g., Pasek and Lauretta, 2008, and references therein) and presumably the
Earth’s core (Scott et al., 2007).
Phosphides are formed during differentiation of rocky planets
and planetesimals with 90–95% of the phosphorus entering the
metal phase to become phosphides (Pasek, 2008). Phosphides are
also formed as the first P mineral condensate in Solar System condensation sequence calculations. Phosphides do not dissolve in
water, and instead react with water to form a variety of P compounds – schreibersite generates phosphate and reduced P oxides
when mixed with water. If organic compounds are dissolved in the
water, phosphonates will also form at abundances proportional to
the amount of organic compounds present (Pasek et al., 2007).
Other phosphides such as AlP and Ca3P2 generate phosphine on
addition of water, however these phosphides are not known to occur naturally.
Phosphine, PH3, is a highly reduced P compound in which the oxidation state is nominally !3. Phosphine is the only volatile natural
P compound. Phosphine has been detected in the atmospheres of
Jupiter (Kunde et al., 1982) and Saturn (Fletcher et al., 2009) where
it is formed as an equilibrium product in H-rich giant planet atmospheres (Visscher et al., 2006). Phosphine also occurs in low abundance in the Earth’s atmosphere (Devai et al., 1988; Glindemann
et al., 1996), formed on the Earth primarily by unidentified biological processes (Devai and Delaune, 1995), and is also generated in
low quantities by a number of high-energy events (Glindemann
et al., 2004). We include in our discussion of phosphine the compounds diphosphine – P2H4 – a highly reactive gas formed by the
photolysis of PH3 (Ferris and Khwaja, 1985), and methylphosphine
– CH3PH2, a gas shown to be produced by phosphates struck by
lightning under reducing conditions (Glindemann et al., 1999).
3. Methods
3.1. Solution chemistry
We investigated the solubility of three classes of P compounds
in three different environments: an organic-rich aqueous environment (Solution 1), an ammonia-rich aqueous environment (Solution 2), and a liquid organic environment (Solution 3). The
organic-rich aqueous environment may be representative of an impact melt lake on Titan’s surface, in which liquid water and watersoluble organics such as acetonitrile interact. The ammonia–water
experiment might be applicable to the liquid layer present beneath
the ice I crust of Titan or in ammonia-rich melts in the ice I crust
(Mitri et al., 2008). The liquid organic environment was chosen
to study the solubility of P compounds in pure hydrocarbons that
would be immiscible in water, perhaps reflective of the partitioning of P compounds as an ice-water melt comes in contact with
non-polar organic compounds.
Of phosphates, we chose orthophosphate as phosphoric acid,
H3PO4, and pyrophosphate as the tetrasodium salt, Na4P2O7. For reduced P oxides, we selected phosphite as phosphorous acid, H3PO3,
hypophosphate as the trisodium salt, Na3HP2O6, and hypophosphite as the sodium salt, NaH2PO2. For phosphonates, we selected
methylphosphonate as methylphosphonic acid, CH3PO3H2, ethylphosphonate as ethylphosphonic acid, CH3CH2PO3H2, isopropylphosphonate as isopropylphosphonic acid, CH3CH2CH2PO3H2,
Author's personal copy
M.A. Pasek et al. / Icarus 212 (2011) 751–761
753
Fig. 1. Phosphorus species discussed in text.
and acetylphosphonate as acetylphosphonic acid, HOOCCH2PO3H2.
These compounds were procured directly from chemical suppliers,
or in the case of hypophosphate were synthesized in house (Pasek
et al., 2007). Concentrations of these compounds in the final, mixed
solution ranged from 0.8 to 1.2 mM (Table 1). The pH of the final
solution was approximately equal to 7–8, based on peak positions
(Yoza et al., 1994).
These compounds were completely soluble in solution, due to
the use of Na salts or their acids, hence all changes that occurred
to the solution upon addition of other compounds were losses
due to insolubility, rather than increases in solubility. We chose
not to investigate the chemistry of phosphine in these systems as
our experimental apparatus was not equipped for handling this
dangerous gas. Phosphides were also not investigated since these
compounds are not soluble in water, but instead release phosphates, reduced P compounds, and phosphonates on reaction with
water; the inclusion of phosphides would hence be redundant.
We diluted 2 mL aliquots of the aqueous prepared solution with
6 mL of either acetonitrile (CH3CN – J.T. Baker, 99.8%, Solution 1),
33% NH4OH in water (VWR, Solution 2), or paraffin oil (VWR,
Author's personal copy
754
M.A. Pasek et al. / Icarus 212 (2011) 751–761
Table 1
P compound abundance in the starting aqueous solution, in
millimolar concentration. All other solutions (1–3) were prepared
using this starting solution.
Compound name
mM
Isopropylphosphonic acid
Ethylphosphonic acid
Methylphosphonic acid
Hypophosphate
Phosphite
Orthophosphate
Pyrophosphate
Acetylphosphonic acid
Hypophosphite
1.084565
0.950235
0.978647
0.786584
1.030363
1.212968
1.145113
0.900714
0.887736
Solution 3). The vials were then shaken to promote mixing of solvents. The temperature of these experiments was 20 "C. A white
precipitate formed in solutions with added acetonitrile and
NH4OH, both of which mixed completely. No precipitate was observed in the water with added oil, which remained separated into
a water layer and oil layer. One mL of each solution was then
mixed with one mL of D2O and was analyzed on a Varian 300 four
nucleus probe FT-NMR spectrometer operating at 121.43 MHz and
24.5 "C for 1024 scans. Spectra were acquired in both H-decoupled
and coupled modes. The abundance of each compound was determined by integration under the peaks of individual species in the
NMR spectrum, following prior work (e.g., Pasek et al., 2007).
Changes in compound abundance directly reflected solubility of P
compounds within solutions with varied solvents, and represent
different affinities of the P compounds for these varied solutions.
While these environments are not exact representations of the
range of environments in which P compounds may be concentrated, they do provide a first-order estimate for some of the processing of P that may occur on outer solar system bodies.
3.2. Thermodynamic equilibrium modeling of phosphine
Since we did not investigate the chemistry of phosphine experimentally, we determined the solubility of phosphine in organicrich water, ammonia-rich water, and in hydrocarbons, to parallel
our experiments. These calculations had as their goal to simulate
the behavior of phosphine in the cold, organic environment of
Titan.
The solubility of PH3 (g) in water has been explored previously
(Schulte et al., 2001). PH3 is a non-electrolyte as it does not readily
react as either an acid or a base. It has a solubility of 0.008 M in
pure water at 298 K and 1 bar of PH3. In solutions of 25% NH3 in
water and 75% acetonitrile in water, we solve for solubility of
PH3 by decreasing the temperature of the reaction:
PH3 ðgÞ ! PH3 ðaqÞ
ðR1Þ
to freezing point at the NH3–H2O eutectic (Kargel et al., 1991;
Mousis et al., 2002) and to a 75:25 CH3CN–H2O mixture, and calculate solubility using Henry’s Law (Wilhelm et al., 1977). These estimations do not include measurements of PH3 solubility in these
mixtures, so at best these estimates give an order of magnitude estimate of solubility. The solubility of PH3 is known to increase in nonpolar solvents, hence our calculated values are likely lower
estimates.
Solubility data for PH3 (g) in linear aliphatic hydrocarbons, and
temperature data for solubility of PH3 in methylbenzene (Young
et al., 1985 and references therein) show that PH3 is most soluble
in larger aliphatic organic compounds (Fig. 2A) and it increases in
solubility as temperatures decrease (Fig. 2B), and is more soluble in
these solvents than in water by a factor of about 100. Using PH3
Fig. 2. Solubility of phosphine in alkanes in terms of ln PH3(dissolved)/PPH3.
PH3(dissolved) is given in units of moles per kg of solvent (errors are smaller than
points). (a) vs. carbon number for linear alkanes at 298 K, with line showing Eq. (1).
(b) Solubility vs. temperature (K) in methylbenzene, with line showing Eq. (2).
solubility data for pentane, hexane, and heptane, we calculated
the best fit solubility relations for the reaction:
PH3 ðgÞ $ PH3 ðsolutionÞ
ðR2Þ
with respect to carbon number at 298 K as:
log Kðin moles PH3 =kg solventÞ ¼ 0:0909C N ! 0:9979 ðR2
¼ 0:9938Þ
ð1Þ
where CN is the number of carbons in the aliphatic compound. Using
data for methylbenzene we calculated solubility with respect to
temperature using the Van’t Hoff equation:
log Kðin moles PH3 =kg solventÞ ¼ 920:61ð1=TÞ ! 3:5445 ðR2
¼ 0:9945Þ
ð2Þ
Note that these are reactions leading from 1 atm of gaseous PH3 to
dissolved PH3. PH3 undergoes a change in state at 185 K and 1 atm,
becoming a liquid. Phosphine’s freezing point is 135 K, where it
forms a-PH3. Below 88 K b-PH3 is the stable phosphine phase, and
c-PH3 below 40 K, and finally d-PH3 at below 30 K (Stephenson
and Giauque, 1937; Hardin and Harvey, 1964). Thus the actual
chemical process that occurs on Titan’s surface is:
PH3 ðsÞ $ PH3 ðsolutionÞ
ðR3Þ
Using HSC Chemistry thermodynamic equilibrium software
(Outokompu Research Oy, see Pasek et al., 2005) we determined
Author's personal copy
M.A. Pasek et al. / Icarus 212 (2011) 751–761
thermodynamic data (DGformation and S) for PH3 (g) at 185 K. Since
at this temperature the DGformation of PH3 (g) is equivalent to PH3
(l), we solved for the entropy using the latent heat of vaporization
and known heat capacities (Stephenson and Giauque, 1937). Similarly we calculated S and DG values for all forms of PH3 to 85 K.
We then calculated the PH3 (s) solubility in organics using the DG
values for solid PH3 – either a or b which at 85 K are very similar
in value, referenced to the difference between these values and
the DG value of PH3 (g). While this is an approximation, it likely
yields a closer estimate of the solubility of PH3 in these fluids at
these temperatures.
On incorporating the phase change from gas to solid, and by
combining the temperature dependence and carbon number
dependence, we solved the solubility of PH3 (s) in ethane and
methane as:
log Kðin mols PH3 =kg C2 H6 Þ ¼ 920:61ð1=TÞ ! 10:139
log Kðin mols PH3 =kg CH4 Þ ¼ 920:61ð1=TÞ ! 10:229
ð3Þ
ð4Þ
4. Results
On addition of acetonitrile to the solution (Solution 1), three
species were completely lost from the solution – acetylphosphonate, hypophosphate, and pyrophosphate (Fig. 3). These compounds had the highest negative charge in this pH range, hence
these compounds have reduced affinity for the poorly polar solvent. In contrast, the reduced P oxides phosphite and hypophosphite were both more soluble in these solutions relative to phosphate,
probably due to the presence of H–P bonds, which share some
affinity with H–C bonds in terms of polarity. The alkyl phosphonates were also more soluble, as the organic groups share similar
chemistry with the organic solution. Acetylphosphonate was less
soluble than these compounds, likely due to its strong polarity.
The addition of ammonia to the solution (Solution 2) resulted in
an increase in relative hypophosphite, pyrophosphate and hypophosphate concentrations, suggesting that these compounds were
more soluble than others due to complexation or stabilization by
NH3. Other compounds, such as the phosphonates and reduced P
oxides, precipitated out in significant quantities. Orthophosphate
and acetylphosphonate were the least soluble in this solution, with
over 75% precipitating out of solution.
The addition of oil to the solution (Solution 3) did not remove
any P compounds from the water. The relative ratios of the P species changed by 5% or less, consistent with the error associated
Fig. 3. Fraction of P compounds left in solution when acetonitrile (Solution 1) or
NH4OH (Solution 2) is added to the starting solution. Concentrations have been
normalized to phosphate. ‘‘PA’’ is phosphonic acid. All values started at close to 1
(see Table 1).
755
with the NMR abundance measurements. No P compounds were
observed in NMR analyses of the oil. For these reasons we propose
that no oxygen-bearing P compound will enter an organic oil-like
phase.
4.1. Phosphine
We calculated the vapor pressure of PH3 at 90 K as
&5 ' 10!7 bar. The corresponding mole fraction of PH3 in liquid
water melts should be &10!10, in acetonitrile–water melts should
be &10!9, and in ammonia–water melts should be &10!8. Phosphine is unlikely to be a major constituent of liquid water at Titan’s
surface.
Intriguingly, phosphine is more soluble in non-polar solvents
like hydrocarbons, and increases solubility as temperature decreases from 298 K. We estimate that at the surface of Titan
(90 K, 1.6 bar pressure), PH3 is soluble to a maximum concentration of &1 mol/kg CH4 or C2H6 in liquid lakes of these two hydrocarbons. This corresponds to mole fractions of 1.6% and 2.9% in
methane and ethane respectively and competes with the highest
solubilities of organics in liquid methane and ethane as calculated
by Cordier et al. (2009). Although HCN is a polar organic compound, it is likely present only at the percent level in these fluids
and hence probably does not affect PH3 solubility.
5. Source materials for P on Titan
We consider two models for the abundance and form of P on Titan. In the Collins et al. (2010) model, Titan’s crust is ancient and
surface features are due primarily to the activity of the hydrocarbon cycle on Titan’s surface. In this case, exogenous material
deposited after Titan’s growth to its present size is the primary
source of P oxides on Titan’s surface, given that within the rocky
core of Titan the majority of P is trapped as phosphates. In an alternative model (Tobie et al., 2006), Titan’s crust is young—less than
or of order of a billion years old—and has thus overturned since Titan formed. In this scenario, cryovolcanism may play a role in
delivering P compounds from the interior of Titan to the surface.
To cover these two extremes, we consider two primary reservoirs
of P on Titan: accretionary material, and exogenous material delivered to the surface of Titan since its differentiation.
5.1. Exogenous delivery of P compounds
Phosphorus is a minor element in meteoritic samples, comprising
between a few tenths of a percent to a few percent of the total weight
of a given meteorite. The iron meteorites have the largest range of
phosphorus abundances from less than 0.01 weight percent in the
IVA to 2 weight percent in the IIG iron meteorites (Willis, 1980). In
contrast, chondrites have fairly regular P abundances of about
800–1000 ppm (Wolf and Palme, 2001; Lodders, 2003).
Phosphorus abundances in the Wild 2 samples were approximately chondritic to within a factor of 2 (Flynn, 2008), consistent
with the majority of meteoritic material. However, the Stardust
samples apparently had experienced a loss of volatiles, suggesting
a dilution of total P abundances by up to a factor of about 2, if the
ice and organic component for comets is 50%.
For the purpose of these calculations we therefore assume that
extraterrestrial impactors hitting Titan had a P abundance that was
half of the chondritic value. We take the relative proportions of P
compounds to be 69.9% phosphate, 0.1% phosphonate, and 30%
phosphide. In contrast, the CM chondrite Murchison has about
99% of its phosphorus in phosphates, about 0.03% in phosphonates
(Cooper et al., 1992), and about 1% as phosphides. However, due to
the detection of schreibersite in Stardust (Leroux et al., 2008), and
Author's personal copy
756
M.A. Pasek et al. / Icarus 212 (2011) 751–761
the implied increase in organics in cometary material over CM
chondrite material (Flynn, 1995), we propose that these numbers
are reasonable estimates of the exogenous P mineral flux to Titan’s
surface. No reduced P oxides or phosphine have been observed in
meteorite samples or comet spectra, suggesting that these compounds evolve through processes in planet-sized bodies, hence
we assume they are not part of the exogenous material that falls
to Titan.
The present cratering rate at Titan’s surface is estimated at one
>10 km diameter crater formed every 5 ' 105–2 ' 106 years
(Zahnle et al., 2003) but with a likely uncertainty of an order of
magnitude (Wood et al., 2010). This is equivalent to a flux of about
2 ' 1015 kg/Gyr at present. Including the LHB, the total material
deposited on Titan’s surface was approximately 1018–1019 kg
(Charnoz et al., 2009). Large impactors make up the majority of
the mass delivered exogenously.
In the Collins et al. (2010) model, this is equivalent to about
1015–1016 kg of P compounds on Titan’s surface, or 107–108
kg/km2, ignoring exogenous dust contributions. In contrast, the
Earth’s crustal P abundance to a depth of 10 km is about 1010
kg/km2. The Tobie et al. (2006) model has significantly less, about
1012 kg of exogenous P compounds on Titan’s surface, or 104
kg/km2, if we assume that the youthful crust did not preserve P
compounds from the late heavy bombardment event.
On the Earth, interplanetary dust particles comprise a major
fraction of the exogenous material delivered to the surface of the
Earth (Love and Brownlee, 1993). While the dust fall at Titan is unknown, using the Earth as a guide one might increase exogenous P
abundances by a factor of 10–100 (e.g., Chyba and Sagan, 1992;
Pasek and Lauretta, 2008).
5.2. Accretionary P compounds
Titan accreted in the saturnian sub-nebula from solids originating from the solar nebula (Mousis et al., 2009a), and formed from
both rocky and icy components in approximately a 1:1 ratio. Titan
is presumed to have formed from planetesimals formed in the solar
nebula that have been thermally processed within Saturn’s subnebula (Mousis et al., 2009b). Only the highly volatile species have
been devolatilized from Titan’s planetesimals during their migration/accretion within Saturn’s sub-nebula (CO, N2, Ar). The other
volatiles remained trapped in planetesimals and were incorporated
into Titan during its formation. Phosphorus in the rocky materials
is taken to be identical to the current exogenous flux, however the
icy phase included substantial amounts of P in the form of PH3
trapped by clathrate hydrates in the solar nebula. It is possible to
quantify the amount of P trapped in the icy part of Titan’s planetesimals by examining the formation sequence of the different ices
produced in the formation zone of Saturn, illustrated in Fig. 4. This
formation sequence is calculated using (i) a predefined initial gas
phase composition of the disk and (ii) the equilibrium curves of
stochiometric hydrates, clathrates and pure condensates, and the
thermodynamic path (hereafter cooling curve) detailing the evolution of temperature and pressure at 9.5 AU roughly corresponding
to the current location of Saturn in the Solar System.
The composition of the initial gas phase of the disk is defined as
follows: we assume that the abundances of all elements considered
(O, C, N, P, S, Ar, Kr and Xe) are protosolar (Asplund et al., 2009) and
that O, C, and N exist only under the form of H2O, CO, CO2, CH3OH,
CH4, N2, and NH3. Concerning the distribution of elements in the
main volatile molecules, we set CO/CO2/CH3OH/CH4 = 70/10/2/1
and N2/NH3 = 1/1 in the gas phase of the disk (values taken from
Mousis et al., 2009b). In addition, S is assumed to exist in the form
of H2S, with H2S/H2 = 0.5 ' (S/H2)solar, and other refractory sulfide
components (Pasek et al., 2005). Once the molecular abundances
are fixed, the remaining O gives the abundance of H2O. The
Fig. 4. Equilibrium curves of NH3–H2O hydrate, H2S, PH3, Xe, CH4 and CO clathrates
(solid lines), CH3OH, CO2, Kr, CO, Ar and N2 pure condensates (dotted lines), and
thermodynamic path followed by the solar nebula at the current position of Saturn
as a function of time, respectively, assuming a full efficiency of clathration. Species
remain in the gas phase above the equilibrium curves. Below, they are trapped as
clathrates or simply condense.
equilibrium curves of hydrates and clathrates derive from Lunine
and Stevenson’s (1985) compilation of published experimental
work, in which data are available at relatively low temperatures
and pressures. The equilibrium curves of pure condensates used
in our calculations derive from the compilation of laboratory data
given in the CRC Handbook of Chemistry and Physics (Lide, 2002).
The thermodynamic path corresponding to the decrease of pressure and temperature conditions of the solar nebula gas at the
heliocentric distance of 9.5 AU intercepts the equilibrium curves
of the different ices at particular temperatures and pressures. For
each ice considered, the domain of stability is the region located
below its corresponding equilibrium curve. The clathration process
stops when no more crystalline water ice is available to trap the
volatile species. Note that, in the pressure conditions of the solar
nebula, because CO2 crystallizes at a higher temperature than its
associated clathrate, we assume that solid CO2 is the only existing
condensed form of CO2 in this environment. In addition, we have
considered only the formation of pure ice of CH3OH in our calculations since, to the best of our knowledge, no experimental data
concerning the equilibrium curve of its associated clathrate have
been reported in the literature.
In the present case, the clathration efficiency is presumed to be
100%, implying that guest molecules had the time to diffuse
through porous water–ice solids before their growth into planetesimals and their accretion by Saturn and its proto-satellites, or that
collisions among planetesimals created fractures and exposed
fresh ice to the gas phase (Lunine and Stevenson, 1985). In this
case, NH3, H2S, PH3, Xe, CH4 and about 30% of available CO form
NH3–H2O hydrate and clathrates a priori dominated by H2S, PH3,
Xe, CH4 and CO with the available water in the outer nebula (updated value from Mousis et al. (2009b)). The remaining CO, as well
as N2, Kr, and Ar, whose clathration normally occurs at lower temperatures, stay in the gas phase until the nebula cools enough to
allow the formation of pure condensates.
In order to estimate the amount of PH3 trapped in clathrates
incorporated in Titan’s building blocks, it is necessary to examine
the relative abundances of guests that can be incorporated in clathrates at the time of their formation in Saturn’s feeding zone. Indeed, the composition of clathrates formed in the solar nebula
was investigated by Mousis et al. (2010) and these authors have
Author's personal copy
M.A. Pasek et al. / Icarus 212 (2011) 751–761
shown that CO becomes a minor compound in the clathrate that in
previous work was expected to be dominated by this molecule. Because the solar composition gas considered by Mousis et al. (2010)
was free of PH3, we reexamine here the trapping conditions of this
molecule in clathrates formed in Saturn’s feeding zone. In our calculations, any volatile already trapped or condensed at a higher
temperature than the formation temperature of the clathrate under consideration is excluded from the coexisting gas phase
composition. This implies that PH3 can only be incorporated in
H2S- and PH3-dominated clathrates.
To calculate the relative abundances of guest species incorporated in a clathrate from a coexisting gas of specified composition
at given temperature and pressure, we follow the method described
by Lunine and Stevenson (1985) and Mousis et al. (2010) using classical statistical mechanics to relate the macroscopic thermodynamic
properties of clathrates to the molecular structure and intermolecular potentials, which depend on parameters describing the interaction between the molecule and the cage. These calculations are
based on the original ideas of van der Waals and Platteeuw (1959)
for clathrate formation, which assume that trapping of guest molecules into cages corresponds to the three-dimensional generalization of ideal localized adsorption. Our calculations have been made
using the PH3 Lennard–Jones parameters derived by Vorotyntsev
and Malyshev (1998) and the set (1) of Kihara parameters used by
Mousis et al. (2010) for the other molecules under consideration.
Set (2) of the Kihara parameters presenter in Mousis et al. (2010)
was incomplete, hence we have used set (1).
Fig. 5 represents the mole fraction of volatiles encaged in structure I and structure II clathrates a priori dominated by H2S (a) and
PH3 (b) and formed in the primordial nebula. This figure shows that
757
PH3 is efficiently trapped in the two clathrates, irrespective of their
structure. Interestingly enough, we note that the mole fraction of
PH3 is larger than the one of H2S in the clathrate expected to be
dominated by this latter molecule. This implies that essentially
all PH3 should be incorporated in the clathrate a priori dominated
by H2S instead of being enclathrated at lower temperature. We
then infer that essentially all PH3 should be incorporated in the
form of clathrates in Titan’s building blocks regardless the clathration efficiency (even a 10% clathration efficiency should be enough
the trap essentially all PH3 in clathrate). Following the approach
presented by Mousis et al. (2009b) for calculating the composition
of ices, we infer that the PH3/H2O ratio is order of 880 ppm in the
icy material accreted in Titan. This value happens to be very close
to the P abundance (about 900 ppm, see Wolf and Palme, 2001)
estimated in the rocky material of Titan, implying that the bulk
composition is order of 890 ppm for an ice-to-rock ratio of one in
the satellite.
The migration of Saturn as a consequence of the dynamical evolution of the Solar System (e.g., Gomes et al., 2005) does not change
these results significantly as Marboeuf et al. (2008) have shown
that the composition of the ices within planetesimals formed in
the outer disk remains almost the same irrespective of their formation distance and of the input parameters of the disk model. Hence,
the migration of Saturn over a reasonable distance (a few AU)
should not affect our conclusions.
6. Fates of P materials
Accretionary and exogenous P compounds were subsequently
modified by a number of processes on Titan’s surface. We consider
here accretionary processing, impact melting, and oxidation and
reduction of P compounds.
6.1. Accretionary processing
The materials forming Titan would have been subjected to
accretionary heating. This likely would not have affected the phosphides or phosphates to any significant degree, as these compounds would be refractory compared to the majority of the
building materials of Titan. However, when the surface temperature has exceeded the melting point of water ice due to accretional
heating, a relatively massive proto-atmosphere was generated
from the release of the volatiles trapped in the icy phase of proto-Titan or in the accreted planetesimals. Assuming that Titan is
roughly 50% ice by mass, the maximum mass of PH3 released into
the atmosphere is about 1.1 ' 1020 kg, which translates into a partial pressure of about 18 bar at the ground level. This value remains
lower than the vapor pressure of 45 bar1 predicted for a surface
temperature of about 300 K after accretion (Kuramoto and Matsui,
1994). However, assuming a PH3 atmospheric pressure similar to
the calculated value, and using the Henry’s constant derived from
the laboratory work of Schulte et al. (2001), the fraction of PH3 dissolved in the ocean was about 0.0025 mol/mole. From these considerations, it appears that significant amounts of PH3 were present in
the primitive ocean of Titan and remained trapped within Titan in
the form of clathrates or pure condensate during its cooling.
6.2. Impact melts
Transient impact melts of liquid water are generated by large
impacts into the icy crust of Titan. These melts can persist for
Fig. 5. Mole fraction of volatiles encaged in H2S (a) and PH3 (b) clathrates. Grey and
dark bars correspond to structure I and structure II clathrates, respectively.
1
Vapor pressure derived from tables located at http://encyclopedia.airliquide.com/
Encyclopedia.asp?GasID=51 and http://encyclopedia.airliquide.com/images_encyclopedie/VaporPressureGraph/Phosphine_Vapor_Pressure.GIF.
Author's personal copy
758
M.A. Pasek et al. / Icarus 212 (2011) 751–761
102–104 years depending on crater size and the makeup of Titan’s
surficial ice (O’Brien et al., 2005; Artemieva and Lunine, 2005;
Neish et al., 2008). Within this time, water will react with impactor
fragments to liberate phosphorus and modify P compound
chemistry.
Schreibersite in an impactor corrodes at a rate of approximately
0.95 g m!2 day!1 under anoxic conditions (Bryant et al., 2009). The
composition of this material would be approximately 30% phosphate, 30% phosphite, 20% hypophosphite, 9% hypophosphate, 9%
pyrophosphate, and 9% phosphonates (Pasek and Lauretta, 2005;
Bryant and Kee, 2006; Pasek et al., 2007). The relative ratios of
these species may vary depending on the corrosion pathway
(e.g., Bryant and Kee, 2006), but are typical products from schreibersite corrosion. Included with these reduced P oxides are phosphonates formed during corrosion, and are based on a corrosion
in organic-rich (&2%) water (Pasek et al., 2007).
As an example, a 15 km impact crater is generated by an impactor of mass &4 ' 1011 kg, and would have about 6 ' 107 kg of
schreibersite. The impactor would sit in a crater filled with meltwater for 102–103 years. The rate of corrosion would be highly
dependent on the fragmentation of the impactor, but as minimum,
we can assume a completely spherical impactor. In this case, about
105–106 kg of schreibersite is corroded within the lifetime of this
impact oasis. Using more reasonable surface areas would result
in total corrosion prior to complete freeze-out of the crater. Larger
craters (150 km) should see complete corrosion due to increased
lifetimes of lakes. Hence nearly all exogenous schreibersite should
form phosphate and other reduced P oxides as a consequence of
reaction with meltwater formed during impact.
As an impact melt freezes out, the composition of the water
may change, as the water is filled with organic compounds like
acetonitrile, or ammonia as the water freezes and approaches the
NH3–H2O eutectic. In these cases, the reduced P oxides phosphite
and hypophosphite, and phosphonates will dominate if organic
compounds start to accumulate with the water, whereas in ammonia-rich water hypophosphate, pyrophosphate, and hypophosphite
will dominate the melt.
6.3. Oxidation
Reduced P oxides formed by corrosion of schreibersite can be further oxidized by reaction with oxidants like H2O2 and O2, forming
phosphates and condensed phosphates (Pasek et al., 2008). However, neither peroxide nor O2 have been observed on Titan. These
compounds may be more important on a planetary body with abundant oxidants like Europa (Moore and Hudson, 2000), as opposed to
Titan, where surface radiation is limited to 4.5 ' 109 eV cm!2 s!1
from cosmic rays (Molina-Cuberos et al., 1999), about 10,000 times
smaller than the surface energy deposition rate for Europa.
Oxidation in anoxic water tends to proceed very slowly for most
reduced P oxides, and it is estimated to take billions of years to oxidize phosphite in water in contact with a reducing atmosphere
(Pasek, 2008). Hence, all reduced P oxides formed by oxidation of
phosphides or by reduction of phosphates should persist for the
crustal lifetime of Titan. The oxidation pathways of phosphine
are unclear, but presumably in the absence of photolysis, it should
persist indefinitely as a solid or solute on Titan’s surface.
6.4. Reduction
High-energy events occurring on the surface of Titan have the
potential to reduce phosphates by reaction with organic material
present in situ at Titan’s surface. Cloud-to-ground lightning has
been reported to reduce phosphate to phosphite (Pasek and Block,
2009). In the presence of a reducing atmosphere, reduction would
take place even more readily (Glindemann et al., 1999). Lightning
has not been detected on Titan (Fischer et al., 2007; Morente
et al., 2008), and convective storms potentially capable of generating lightning will be much less abundant on Titan than on the Earth
due to the lower solar flux at the former (Awal and Lunine, 1994).
7. Transport of P materials
Modification of P compounds on the surface of Titan may result
in a diversity of P compounds that differ in their mobility in the
variety of fluids that are plausible on Titan’s surface.
Titan’s internal structure is a core composed of a rocky material
of chondritic composition, surrounded by ices with a liquid water
or water–ammonia layer sandwiched between the high and low
pressure ices, with a possible mixed rock–ice layer in place of or
part of the high pressure ice layer (Iess et al., 2010). This separation
of the core and liquid subsurface by high pressure ices suggests
that refractory material trapped in the core (e.g., phosphides and
phosphates) could not be released to the surface of Titan via
cryovolcanism.
The transport of phosphates, reduced P compounds, and phosphonates will be highly contingent on meltwater composition.
Meltwaters rich in ammonia will carry primarily hypophosphite,
hypophosphate, and pyrophosphate from exogenous sources,
whereas meltwaters rich in organics will transport phosphonates,
phosphite, and hypophosphite. The extent to which this process
might be important on the surface of Titan cannot be quantified
with the existing Cassini data.
In contrast to its effects on potassium (Engel et al., 1994),
ammonia does not increase the solubility of phosphine, as ammonia and phosphine do not share very many chemical characteristics. Liquid water on the surface of Titan, generated by either
cryovolcanism or impacts, would likely carry phosphine at about
a concentration of &10!10 mol/mole, as its solubility is limited by
the vapor pressure of PH3 at Titan’s surface (see reaction R1).
The major difference between P chemistry on Earth and on Titan results from the chemistry of phosphine. Phosphine is more
soluble in organic compounds than it is in water. Hence, PH3 will
participate in the hydrocarbon cycle. Hydrocarbons like methane
and ethane which are stable liquids at Titan’s surface will transport
PH3. Furthermore, these hydrocarbons may also preferentially extract PH3 from crustal ice rock through mechanical and chemical
erosion. Because the hydrocarbon liquids in which PH3 dissolves
undergo seasonal and longer-term evaporation/condensation cycles (Aharonson et al., 2009), evaporites containing phosphine
could be present around the lakes and in dry lakebeds, and might
be detected with future remote sensing instruments in Titan orbit.
8. Discussion
Phosphine likely controls P chemistry on satellites at the orbit of
Saturn and beyond. At Jupiter and inward, phosphates likely control
P chemistry. Phosphorus chemistry is likely linked to the snow line
(e.g., Lunine et al., 2000), with P chemistry inward of the snow line
dominated by phosphates, and outward dominated by PH3.
If phosphine dominates P chemistry on Titan’s surface, then it
merits a brief discussion of possible reactions phosphine can undergo. If an H–P bond is broken either by photolysis by UV light
or by spark discharges then the PH2 radical can participate in several reactions (Ferris et al., 1984; Ferris and Khwaja, 1985; Bossard
et al., 1986; Guillemin et al., 1995). These reactions include formation of diphosphine:
2PH3 ! P2 H4 þ H2
which in turn photolyzes to red phosphorus:
P2 H4 ! P ðsÞ þ H2
Author's personal copy
M.A. Pasek et al. / Icarus 212 (2011) 751–761
Reactions with atmospheric organic constituents are promoted by
spark discharges (Bossard et al., 1986).
PH3 þ CH4 ! CH3 PH2 þ H2
These reactions include reaction with acetylene to form ethynylphosphine via UV light:
PH3 þ C2 H2 ! HC2 PH2 þ H2
In the absence of high energy UV light (which is essentially fully absorbed in the atmosphere well above the surface) and without
cloud-to-ground lightning (a phenomenon for which there is no evidence but, if present on Titan, would supply 0.1–1 GJ per strike, see
Krider et al., 1968), phosphine would be stable on Titan’s surface
and would likely not form these compounds extensively. However,
if PH3 enters the upper atmosphere, perhaps if volatilized by a large
impact, then PH3 could react with atmospheric methane, generating
organophosphorus compounds like methylphosphine. Furthermore,
cosmic rays generate ionizing radiation and deposit kinetic energy
directly into the surface at a total estimated rate of 109 eV cm!2 s!1,
enough to produce 3 ' 10!17 g cm!2 s!1 of benzene from acetylene
(Zhou et al., 2010). Some of this energy will go into processing of
phosphine and reactions of phosphine with adjacent organic
molecules.
The formation of organophosphine compounds like methyl- and
ethylphosphine proceeds most readily by spark discharges, which
may model approximately the effect of electrons produced by cosmic rays. Other compounds may form if PH3 reacts with trace
organics like ethane or acetylene, although these will tend to form
at the abundance of these individual gases and will only substitute
for one of the three hydrogens of PH3 (Bossard et al., 1986). However, if ethenylphosphine is formed to any significant concentration (Guillemin et al., 1995), then cosmic ray processing of this
compound may lead to phenylphosphine, in analogy to the processes that lead to benzene from acetylene (Zhou et al., 2010). In
this respect, C6H5–PH2 and phosphorylated organics in general
may comprise a small portion of the haze of Titan.
Phosphorus is sensitive to detection by Nuclear Magnetic Resonance Spectroscopy. Phosphorus oxide compounds are characterized by spectral peaks occurring in the !30 to 35 ppm range
(referenced to 85% H3PO4). Phosphine and associated compounds
(e.g., methylphosphine, phenylphosphine, and others) in contrast
occur at the !120 to !250 ppm range. Additionally, these compounds may be strong and distinctive enough to be detected in
1
H spectra.
Phosphate and organic phosphate compounds on Titan are
likely confined to transient impact crater melts or cryovolcanic
flows, as these compounds do not form readily or react in organic
solvents. In this respect, ‘‘Life as we know it’’ would be confined to
activity during stochastic impacts producing meltwater, and would
have to lie dormant after the surface of Titan froze over these lakes.
Similarly, if life is dependent on phosphate biomolecules, then it
would have had to have developed relatively quickly on Titan’s
surface, on the order of 104–105 years (the maximum longevity
of impact lakes – see O’Brien et al., 2005). The highly soluble hypophosphite appears to be the major P compound capable of participating in chemistry in meltwaters of varied composition. Recent
investigations of hypophosphite chemistry (e.g., Bryant et al.,
2010) have shown this compound to be quite versatile, hence hypophosphite-based life is the most likely of all P–O molecular systems to occur on Titan.
Alternatively, ‘‘life as we don’t know it’’ may incorporate phosphine in its biochemistry on Titan. Phosphine on Titan readily migrates as a result of increased solubility in organic solvents, hence
PH3 would likely not be a limiting reagent, unlike geobiological
systems on the Earth (e.g., Benitez-Nelson, 2000). Additionally, organic compounds readily substitute for H on PH3, hence P may still
759
be relevant in any putative biochemical systems that may live on
Titan. Indeed, phosphine has been suggested as a possible alternative biomolecule with a dipole in Titan’s lakes (Naganuma and
Sekine, 2010), and suggests a greater role for this compound in
alternative life than has been previously suggested.
9. Conclusion
We have determined the P compounds likely present on Titan’s
surface, the chemical modifications of P that may take place, and
the transport pathways of these compounds. Phosphorus chemistry on Titan varies substantially from P chemistry on the Earth
and terrestrial planets. Titan’s surface P chemistry is dominated
by exogenous P compounds like phosphates, reduced oxidation
state P compounds, and phosphonates, and by accretionary phosphine. Accretionary phosphine arrived in clathrates during Titan’s
formation, and was concentrated in the crust following the freezing of the primordial ocean. Exogenous P compounds excluding
phosphine are refractory on Titan’s surface, and will only be transported or change during transient heating events at Titan’s surface.
In contrast, Titan’s active P chemistry is controlled by phosphine.
Phosphine is capable of participating in the hydrocarbon cycle of
Titan, as it is highly soluble in organic solvents. Phosphine may also
participate in haze chemistry in Titan’s atmosphere, forming more
complex species like diphosphine and organic phosphines by photolysis and other high-energy events on Titan. The dominance of
phosphine and the poor mobility of exogenous phosphates and
other P oxides suggests that life on Titan, if present, would either
be free of P or would use phosphine in place of phosphates in its
biochemistry.
Acknowledgments
The authors thank V. Pasek for assistance with figures. This
work was supported by a grant from NASA Exobiology and Evolutionary Biology NNX07AU08G.
References
Abbas, O., Schulze-Makuch, D., 2002. Acetylene-based pathways for prebiotic
evolution on Titan. Int. J. Astrobiol. 1, 233 (abstract).
Aharonson, O., Hayes, A.G., Lunine, J.I., Lorenz, R.D., Allison, M.D., Elachi, C., 2009. An
asymmetric distribution of lakes on Titan as a possible consequence of orbital
forcing. Nat. Geosci. 2, 851–854.
Artemieva, N., Lunine, J.I., 2005. Impact cratering on Titan: II. Global melt, escaping
ejecta and aqueous alteration of surface organics. Icarus 175, 522–533.
Asplund, M., Grevesse, N., Sauval, A.J., Scott, P., 2009. The chemical composition of
the Sun. Annu. Rev. Astron. Astrophys. 47, 481–522.
Awal, M., Lunine, J.I., 1994. Moist convective clouds in Titan’s atmosphere. Geophys.
Res. Lett. 21, 2491–2494.
Benitez-Nelson, C.R., 2000. The biogeochemical cycling of phosphorus in marine
systems. Earth Sci. Rev. 51, 109–135.
Benner, S.A., Ricardo, A., Carrigan, M.A., 2004. Is there a common chemical model for
life in the universe? Curr. Opin. Chem. Biol. 8, 672–689.
Bossard, A.R., Kamga, R., Raulin, F., 1986. Gas phase synthesis of organophosphorus
compounds and the atmosphere of the giant planets. Icarus 67, 305–324.
Bryant, D.E., Kee, T.P., 2006. Direct evidence for the availability of reactive, water
soluble phosphorus on the early Earth. H-Phosphinic acid from the Nantan
meteorite. Chem. Commun. 22, 2344–2346.
Bryant, D.E., Greenfield, D., Walshaw, R.D., Evans, S.M., Nimmo, A.E., Smith, C.,
Wang, L., Pasek, M.A., Kee, T.P., 2009. Electrochemical studies of iron meteorites.
Phosphorus redox chemistry on the early Earth. Int. J. Astrobiol. 8, 27–36.
Bryant, D.E., Marriott, K.E.R., Macgregor, S.A., Fishwick, C.W.G., Pasek, M.A., Kee, T.P.,
2010. Plausible prebiotic ancestors of sugar–phosphates. Chem. Commun. 46,
3726–3728.
Charnoz, S., Morbidelli, A., Dones, L., Salmon, J., 2009. Did Saturn’s rings form during
the Late Heavy Bombardment? Icarus 199, 413–428.
Chyba, C.F., Sagan, C., 1992. Endogenous production, exogenous delivery, and
impact-shock synthesis of organic molecules—An inventory for the origins of
life. Nature 355, 125–132.
Committee on the Origin, Evolution of Life, 2007. The Limits of Organic Life in
Planetary Systems. National Academies Press. 116pp.
Author's personal copy
760
M.A. Pasek et al. / Icarus 212 (2011) 751–761
Collins, G.C., McKinnon, W.B., Moore, J.M., Nimmo, F., Pappalardo, R.T., Prockter,
L.M., Schenk, P.M., 2010. Tectonics of the outer planet satellites. In: Watters, T.,
Schultz, R. (Eds.), Planetary Tectonics. Cambridge Univ. Press, pp. 264–350.
Cooper, G.W., Onwo, W.M., Cronin, J.R., 1992. Alkylphosphonic-acids and
sulfonic-acids in the Murchison meteorite. Geochim. Cosmochim. Acta 56,
4109–4115.
Cordier, D., Mousis, O., Lunine, J.I., Lavvas, P., Vuitton, V., 2009. An estimate of the
chemical composition of Titan’s lakes. Astrophys. J. 707, L128–L131.
Costanzo, G., Saladino, R., Crestini, C., Ciciriello, F., Di Mauro, E., 2007. Nucleoside
phosphorylation by phosphate minerals. J. Biol. Chem. 282, 16729–16735.
de Graaf, R.M., Visscher, J., Schwartz, A.W., 1995. A plausibly prebiotic synthesis of
phosphonic acids. Nature 378, 474–477.
Devai, I., Delaune, R.D., 1995. Evidence for phosphine production and emission from
Louisiana and Florida marsh soils. Org. Geochem. 23, 277–279.
Devai, I., Felfoldy, L., Wittner, I., Plosz, S., 1988. Detection of phosphine: New aspects
of the phosphorus cycle in the hydrosphere. Nature 333, 343–345.
Dyhrman, S.T., Benitez-Nelson, C.R., Orchard, E.D., Haley, S.T., Pellechia, P.J., 2009. A
microbial source of phosphonates in oligotrophic marine systems. Nat. Geosci.
2, 696–699.
Engel, S., Lunine, J.I., Norton, D.L., 1994. Silicate interactions with ammonia–water
fluids on early Titan. J. Geophys. Res. 99, 3745–3752.
Ferris, J.P., Khwaja, H., 1985. Laboratory simulation of PH3 photolysis in the
atmospheres of Jupiter and Saturn. Icarus 62, 415–424.
Ferris, J.P., Bossard, A., Khwaja, H., 1984. Mechanism of phosphine photolysis.
Application to jovian atmospheric photochemistry. J. Am. Chem. Soc. 106, 318–
324.
Fischer, G., Gurnett, D.A., Kurth, W.S., Farrell, W.M., Kaiser, M.L., Zarka, P., 2007.
Nondetection of Titan lightning radio emissions with Cassini/RPWS after 35
close Titan flybys. Geophys. Res. Lett. 34, L22104.
Fletcher, L.N., Orton, G.S., Teanby, N.A., Irwin, P.G.J., 2009. Phosphine on Jupiter and
Saturn from Cassini/CIRS. Icarus 202, 543–564.
Flynn, G.J., 1995. The delivery of organic matter from asteroids and comets to the
early surface of Mars. Earth Moon Planets 71, 469–474.
Flynn, G.J., 2008. Physical, chemical, and mineralogical properties of Comet 81P/
Wild 2 particles collected by Stardust. Earth Moon Planets 102, 447–459.
Glindemann, D., Bergmann, A., Stottmeister, U., Gassmann, G., 1996. Phosphine in
the lower terrestrial troposphere. Naturwissenschaften 83, 131–133.
Glindemann, D., de Graaf, R.M., Schwartz, A.W., 1999. Chemical reduction of
phosphate on the primitive Earth. Orig. Life Evol. B 29, 555–561.
Glindemann, D., Edwards, M., Schrems, O., 2004. Phosphine and methylphosphine
production by simulated lightning – A study for the volatile phosphorus cycle
and cloud formation in the Earth atmosphere. Atmos. Environ. 38, 6867–6874.
Gomes, R., Levison, H.F., Tsiganis, K., Morbidelli, A., 2005. Origin of the cataclysmic
late heavy bombardment period of the terrestrial planets. Nature 435, 466–469.
Gorrell, I.B., Wang, L., Marks, A.J., Bryant, D.E., Bouillot, F., Goddard, A., Heard, D.E.,
Kee, T.P., 2006. On the origin of the Murchison meteorite phosphonates.
Implications for pre-biotic chemistry. Chem. Comm. 2006, 1643–1645.
Guillemin, J.-C., Janati, T., Lassalle, L., 1995. Photolysis of phosphine in the presence of
acetylene and propyne, gas mixtures of planetary interest. Adv. Space Res. 16, 85–
92.
Hardin, A.H., Harvey, K.B., 1964. Infrared absorption of solid phosphine. Can. J.
Chem. 42, 84–89.
Iess, L., Rappaport, N.J., Jacobson, R.A., Racioppa, P., Stevenson, D.J., Tortora, P.,
Armstrong, J.W., Asmar, S.W., 2010. Gravity field, shape and moment of inertia
of Titan. Science 327, 1367–1369.
Kakegawa, T., Noda, M., Nannri, H., 2002. Geochemical cycles of bio-essential
elements on the early Earth and their relationships to the origin of life. Resource
Geol. 52, 83–89.
Kargel, J.S., Croft, S.K., Lunine, J.I., Lewis, J.S., 1991. Rheological properties of
ammonia–water liquids and crystal–liquid slurries: Planetological applications.
Icarus 89, 93–112.
Keefe, A.D., Miller, S.L., 1995. Are polyphosphates or phosphate esters pre-biotic
reagents? J. Mol. Evol. 41, 693–702.
Kornberg, A., 1950. Reversible enzymatic synthesis of diphosphopyridine nucleotide
and inorganic pyrophosphate. J. Biol. Chem. 182, 779–793.
Krider, E.P., Dawson, G.A., Uman, M.A., 1968. Peak power and energy dissipation in a
single-stroke lightning flash. J. Geophys. Res. 73, 3335–3339.
Kunde, V., Hanel, R., Maguire, W., 1982. The tropospheric gas composition of
Jupiter’s north equatorial belt (NH3, PH3, CH3 D, GeH4, H2O) and the jovian D/H
isotopic ratio. Astrophys. J. 263, 443–467.
Kuramoto, K., Matsui, T., 1994. Formation of a hot proto-atmosphere on the
accreting giant icy satellite: Implications for the origin and evolution of Titan,
Ganymede, and Callisto. J. Geophys. Res. 99, 21183.
LeRoux, H. et al., 2008. A TEM study of thermally modified Comet 81P/Wild 2 dust
particles by interactions with the aerogel matrix during the Stardust capture
process. Meteorit. Planet. Sci. 43, 97–120.
Lide, D.R., 2002. CRC Handbook of Chemistry and Physics: A Ready-Reference Book
of Chemical and Physical Data, 83rd ed. CRC Press, Boca Raton. ISBN
0849304830.
Lodders, K., 2003. Solar system abundances and condensation temperatures of the
elements. Astrophys. J. 591, 1220–1247.
Love, S.G., Brownlee, D.E., 1993. A direct measurement of the terrestrial mass
accretion rate of cosmic dust. Science 262, 550–553.
Lunine, J.I., Stevenson, D.J., 1985. Thermodynamics of clathrate hydrate at low and
high pressures with application to the outer Solar System. Astrophys. J. Suppl.
Ser. 58, 493–531.
Lunine, J.I., Owen, T.C., Brown, R.H., 2000. The outer Solar System: Chemical
constraints at low temperatures on planet formation. In: Mannings, V., Boss,
A.P., Russel, S. (Eds.), Protostars and Planets IV. Univ. of Arizona Press, Tucson,
pp. 1055–1080.
Marboeuf, U., Mousis, O., Ehrenreich, D., Alibert, Y., Cassan, A., Wakelam, V.,
Beaulieu, J.-P., 2008. Composition of ices in low mass extrasolar planets.
Astrophys. J. 681, 1624–1630.
McKay, C.P., Smith, H., 2005. Possibilities for methanogenic life in liquid methane on
the surface of Titan. Icarus 178, 274–276.
Mitri, G., Showman, A.P., Lunine, J.I., Loes, R.M.C., 2008. Resurfacing of Titan by
ammonia–water cryovolcanism. Icarus 196, 216–224.
Molina-Cuberos, G.J., López-Morena, J.J., Rodrigo, R., Lara, L.M., O’Brien, K., 1999.
Ionization by cosmic rays of the atmosphere of Titan. Planet. Space Sci. 47,
1347–1354.
Moore, M.H., Hudson, R.L., 2000. IR detection of H2O2 at 80 K in ion-irradiated
laboratory ices relevant to Europa. Icarus 145, 282–288.
Morente, J.A., Porti, J.A., Salinas, A., Navarro, E.A., 2008. Evidence of electrical activity
on Titan drawn from the Schumann resonances sent by Huygens probe. Icarus
195, 802–811.
Mousis, O., Pargamin, J., Grasset, O., Sotin, C., 2002. Experiments in the NH3–H2O
system in the [0, 1 GPa] pressure range – Implications for the deep liquid layer
of large icy satellites. Geophys. Res. Lett. 29, 20192, doi:10.1029/
2002GL015812.
Mousis, O., Marboeuf, U., Lunine, J.I., Alibert, Y., Fletcher, L.N., Orton, G.S., Pauzat, F.,
Ellinger, Y., 2009a. Determination of the minimum masses of heavy elements in
the envelopes of Jupiter and Saturn. Astrophys. J. 696, 1348–1354.
Mousis, O. et al., 2009b. Clathration of volatiles in the solar nebula and implications
for the origin of Titan’s atmosphere. Astrophys. J. 691, 1780–1786.
Mousis, O., Lunine, J.I., Picaud, S., Cordier, D., 2010. Volatile inventories in clathrate
hydrates formed in the primordial nebula. Faraday Discuss. 147, 509–525.
Naganuma, T., Sekine, Y., 2010. Hydrocarbon lakes and watery matrices/habitats for
life on Titan. J. Cosmol. 5, 905–911.
Neish, C.D., Somogyi, A., Imanaka, H., Lunine, J.I., Smith, M.A., 2008. Rate
measurements of the hydrolysis of complex organic macromolecules in cold
aqueous solutions: Implications for prebiotic chemistry on the early Earth and
Titan. Astrobiology 8, 273–287.
Neish, C.D., Somogyi, A., Smith, M.A., 2010. Titan’s primordial soup: Formation of
amino acids via low-temperature hydrolysis of tholins. Astrobiology 10, 337–
347.
O’Brien, D.P., Lorenz, R.D., Lunine, J.I., 2005. Numerical calculations of the longevity
of impact oases on Titan. Icarus 173, 243–253.
Olsen, E., Fredriksson, K., 1966. Phosphates in iron and pallasite meteorites.
Geochim. Cosmochim. Acta 30, 461–470.
Pasek, M.A., 2008. Rethinking early Earth phosphorus geochemistry. Proc. Natl.
Acad. Sci. USA 105, 853–858.
Pasek, M.A., Block, K., 2009. Lightning reduction of phosphate: Implications for
phosphorus biogeochemistry. Nat. Geosci. 2, 553–556.
Pasek, M.A., Lauretta, D.S., 2005. Aqueous corrosion of phosphide minerals from iron
meteorites: A highly reactive source of prebiotic phosphorus on the surface of
the early Earth. Astrobiology 5, 515–535.
Pasek, M.A., Lauretta, D.S., 2008. Extraterrestrial flux of potentially prebiotic C, N,
and P to the early Earth. Orig. Life Evol. B 38, 5–21.
Pasek, M.A., Milsom, J.A., Ciesla, F.J., Lauretta, D.S., Sharp, C., Lunine, D.S., 2005.
Sulfur chemistry in protoplanetary nebulae with time-varying oxygen
abundances. Icarus 175, 1–14.
Pasek, M.A., Dworkin, J.P., Lauretta, D.S., 2007. A radical pathway for organic
phosphorylation during schreibersite corrosion with implications for the origin
of life. Geochim. Cosmochim. Acta 71, 1721–1736.
Pasek, M.A., Kee, T.P., Bryant, D.E., Pavlov, A.A., Lunine, J.I., 2008. Production of
potentially prebiotic condensed phosphates by phosphorus redox chemistry.
Angew. Chem. Intl. Ed. 47, 7918–7920.
Pech, H., Henry, A., Khachikian, C.S., Salmassi, T.M., Hanrahan, G., Foster, K.L., 2009.
Detection of geothermal phosphite using high-performance liquid
chromatography. Environ. Sci. Technol. 43, 7671–7675.
Pilling, S., Andrade, D.P., Neto, A.C., Rittner, R., Naves de Brito, A., 2009. DNA
nucleobase synthesis at Titan atmosphere analog by soft X-rays. J. Phys. Chem.
113, 11161–11166.
Saladino, R., Crestini, C., Ciciriello, F., Costanzo, G., Di Mauro, E., 2006. Syntheses of
nucleobases and favourable thermodynamic niches for early polymers. Orig.
Life Evol. B. 36, 523–531.
Schulte, M.D., Shock, E.L., Wood, R.H., 2001. The temperature dependence of the
standard-state thermodynamic properties of aqueous nonelectrolytes.
Geochim. Cosmochim. Acta 65, 3919–3930.
Schulze-Makuch, D., Grinspoon, D.H., 2005. Biologically enhanced energy and
carbon cycling on Titan? Astrobiology 5, 560–567.
Schwartz, A.W., 1971. Phosphate: Solubilization and activation on the primitive
Earth. In: Buvet, R., Ponnamperuma, C. (Eds.), Chemical Evolution and the Origin
of Life. American Elsevier Publishing Co., New York, pp. 207–215.
Scott, H.P., Huggins, S., Frank, M.R., Maglio, S.J., Martin, C.D., Meng, Y., Santillan, J.,
Williams, Q., 2007. Equation of state and high-pressure stability of Fe3Pschreibersite: Implications for phosphorus storage in planetary cores. Geophys.
Res. Lett. 34, L06302, doi:10.1029/2006GL029160.
Shapiro, R., Schulze-Makuch, D., 2009. The search for alien life in our Solar System:
Strategies and priorities. Astrobiology 9, 1–9.
Stephenson, C.C., Giauque, W.F., 1937. A test of the third law of thermodynamics by
means of two crystalline forms of phosphine: The heat capacity, heat of
Author's personal copy
M.A. Pasek et al. / Icarus 212 (2011) 751–761
vaporization, and vapor pressure of phosphine. Entropy of the gas. J. Chem.
Phys. 5, 149–158.
Tobie, G., Lunine, J.I., Sotin, C., 2006. Episodic outgassing as the origin of
atmospheric methane on Titan. Nature 440, 61–64.
van der Waals, J.H., Platteeuw, J.C., 1959. Clathrate solutions. Advances in Chemical
Physics, vol. 2. Interscience, New York, pp. 1–57.
Visscher, C., Lodders, K., Fegley Jr., B., 2006. Atmospheric chemistry in giant planets,
brown dwarfs, and low-mass dwarf stars: III. Sulfur and phosphorus. Astrophys.
J. 648, 1181–1195.
Vorotyntsev, V.M., Malyshev, V.M., 1998. Gas hydrates, a new class of impurities in
high purity gases and gas–vapour mixtures. Russ. Chem. Rev. 67, 81–92.
Wilhelm, E., Battino, R., Wilcock, R.J., 1977. Low pressure solubility of gases in liquid
water. Chem. Rev. 77, 219–262.
Willis, J., 1980. The Bulk Composition of Iron Meteorite Parent Bodies. Dissertation,
University of California at Los Angeles.
761
Wolf, D., Palme, H., 2001. The Solar System abundances of phosphorus and titanium
and the nebular volatility of phosphorus. Meteorit. Planet. Sci. 36, 559–571.
Wood, C.A., Lorenz, R.D., Kirk, R., Lopes, R., Mitchell, K., Stofan, E., and the Cassini
RADAR Team, 2010. Impact craters on Titan. Icarus 206, 334–344.
Young, C.L., Fogg, P.G.T., Clever, H.L., Hayduk, W., 1985. Ammonia, Amines,
Phosphine, Arsine, Stibine, Silane, Germane and Stannane in Organic Solvents.
Pergamon Press, Oxford. 360pp.
Yoza, N., Ueda, N., Nakashima, S., 1994. pH-dependence of P-31-NMR spectroscopic
parameters of monofluorophosphate, phosphate, hypophosphate, phosphonate,
phosphinate and their dimers and trimers. Fresen. J. Anal. Chem. 348, 633–638.
Zahnle, K., Schenk, P., Levison, H., Dones, L., 2003. Cratering rates in the outer Solar
System. Icarus 163, 263–289.
Zhou, L., Zheng, W., Kaiser, R.I., Landera, A., Mebel, A.M., Liang, M.-C., Yung, Y.L.,
2010. Cosmic-ray-mediated formation of benzene on the surface of Saturn’s
moon Titan. Astrophys. J. 718, 1243–1251.