Draft version October 1, 2019Typeset using LATEX twocolumn style in AASTeX62

Brown Dwarf Atmospheres As The Potentially Most Detectable And Abundant Sites For Life

Manasvi Lingam1, 2 and Abraham Loeb1

1Institute for Theory and Computation, Harvard University, Cambridge MA 02138, USA2Department of Aerospace, Physics and Space Sciences, Florida Institute of Technology, Melbourne FL 32901, USA

ABSTRACT

We show that the total habitable volume in the atmospheres of cool brown dwarfs with effectivetemperatures of ∼ 250-350 K is possibly larger by two orders of magnitude than that of Earth-likeplanets. We also study the role of aerosols, nutrients and photosynthesis in facilitating life in browndwarf atmospheres. Our predictions might be testable through searches for spectral edges in thenear-infrared and chemical disequilibrium in the atmospheres of nearby brown dwarfs that are eitherfree-floating or within several AU of stars. For the latter category, we find that the James WebbSpace Telescope (JWST) may be able to achieve a signal-to-noise ratio of ∼ 5 after a few hoursof integration time per source for the detection of biogenic spectral features in ∼ 103 cool brown dwarfs.

1. INTRODUCTION

Whenever environments conducive to the origin andsustenance of life (i.e., habitable environments) are stud-ied, they are almost invariably assumed to be based ei-ther on the surface or beneath the surface. The Earthis a classic example of the former, whereas the subsur-face oceans of Europa and Enceladus constitute well-known candidates for the latter (Shapley 1967; Nimmo& Pappalardo 2016). However, there is one crucial en-vironment that is often neglected in considerations ofhabitability, namely, the atmosphere (Schulze-Makuch& Irwin 2008; Preston & Dartnell 2014). The absenceof a continuous solid substrate and reduced protectionagainst cosmic rays as well as ultraviolet (UV) radia-tion constitute some of the reasons commonly advancedto justify why environmental conditions are not suitablefor the presence of aerial biospheres.

However, at an appropriate altitude, favorable physic-ochemical conditions may exist for life, including liquidwater, moderate temperatures and nutrients. Severalstudies have proposed that the cloud layer of Venus at∼ 50 km is potentially habitable in the above respects,even though its surface is far too hot for life-as-we-know-it. This proposal dates back to more than 50 years ago(Morowitz & Sagan 1967) (see also Seckbach & Libby1970) and has gained some traction from the 1990s on-wards (Grinspoon 1997; Cockell 1999; Schulze-Makuchet al. 2004; Dartnell et al. 2015; Limaye et al. 2018).

Apart from Venus, a few early studies also exam-ined the possibility of life in the atmospheres of Jo-vian planets (Sodek & Redmond 1967; Shapley 1967;

Corresponding author: Manasvi Lingam

manasvi.lingam@cfa.harvard.edu

Sagan & Salpeter 1976). These papers were comple-mented by laboratory experiments mimicking Jovian at-mosphere that yielded a number of valuable prebioticcompounds such as amino and imino acids (after acidhydrolysis), aminonitriles, formaldehyde and hydrogencyanide (Sagan 1960; Ponnamperuma & Molton 1973;Ponnamperuma 1976; Stribling & Miller 1987); numer-ical simulations also indicate that some of the abovemolecules as well as considerable amounts of small hy-drocarbons could exist in the upper atmospheres of sub-stellar objects (Bilger et al. 2013; Stark et al. 2014).

The dividing line between giant planets and browndwarfs is not a very sharply delineated one, as therehas been much debate regarding the classification of ob-jects with masses of ∼ 3-10MJ , where MJ denotes themass of Jupiter, and with effective temperatures < 500K; these objects do not appear to possess the capacityfor deuterium fusion (Caballero 2018). As some of thesalient physical and chemical characteristics are similaracross giant planets and cool brown dwarfs (Burrowset al. 2001; Hubbard et al. 2002; Helling & Casewell2014; Bailey 2014; Marley & Robinson 2015), it is natu-ral to inquire whether the atmospheres of brown dwarfsare conducive to the origin and sustenance of life.

To the best of our knowledge, the only paper thathas addressed the issue of atmospheric habitability forcool brown dwarfs (of spectral class Y) is Yates et al.(2017).1 By drawing upon an organism lifecycle model,Yates et al. (2017) found that putative microbes up toan order of magnitude larger, and masses a few orders ofmagnitude higher, than typical microbes on Earth could

1 However, an early qualitative exploration of potential lifeformsin the atmospheres of brown dwarfs and gas giants can be foundin the imaginative, but sadly forgotten, book by Shapley (1967).

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be supported in the presence of atmospheric convection.In that paper, the number of free-floatng cool Y dwarfsin the Miky Way was also investigated, and an estimateof ∼ 109 was derived based on observational constraints.Our work will deal with other aspects of brown dwarfatmospheric habitability and putative biospheres thatwere not explicitly tackled in Yates et al. (2017).

The outline of our paper is as follows. In Section 2,we estimate the maximum habitable volume that is en-capsulated by the atmospheres of cool brown dwarfs.We continue by exploring the characteristics of putativeaerial biospheres in Section 3 with a focus on prebioticchemistry, abiogenesis, and nutrient and energy sup-plies. Subsequently, we identify possible biosignaturesthat might arise in these atmospheres and the methodsof detecting them in Section 4. Finally, we summarizeour central results in Section 5.

2. ASSESSING THE MAXIMUM HABITABLEVOLUME

We shall start by presenting heuristic estimates for thetotal potentially habitable volume encompasses by twodifferent classes of objects. Strictly speaking, our analy-sis yields a potential upper bound for the habitable vol-ume if habitability is narrowly interpreted as conditionspermitting the existence of liquid water. Although we donot explicitly employ the word “potential” henceforth,it should be understood that we are dealing with: (i) po-tential habitability and, (ii) a potential upper bound onthe habitable volume. The “volume” under question isfour-dimensional as it encompasses not only the spatialcoverage but also the temporal duration of habitability.

For a single object, we denote the spatial habitablevolume by V and the habitability interval by τ , yieldingthe four-dimensional (4D) volume (V) defined as V ∼V τ . However, this applies only to a single object. Inorder to determine the total 4D volume spanned by agiven class of objects, we introduce the function

Γ =

∫V dN =

∫V τ dN, (1)

where dN is the number of objects within a particu-lar interval, as we shall describe hereafter. One of thechief points worth appreciating is that V , τ and dN arefunctions of the object’s mass.

2.1. Atmospheric habitable zone in brown dwarfs

The first class of systems we evaluate are the the atmo-spheric habitable regions of brown dwarfs. We use thesubscript ‘BD’ subsequently to identify this category.

We begin by evaluating τBD. This is found by calculat-ing the duration of time over which the brown dwarf hasan effective temperature of ∼ 250-350 K. There are threereasons behind our choice of this thermal range: (a) anumber of interesting molecular species from the per-spective of prebiotic chemistry and biochemistry may be

present, (b) the existence of aerosols is feasible, and (c)atmospheric water vapor and clouds comprising volatileices are believed to exist (Morley et al. 2014; Skemeret al. 2016; Morley et al. 2018; Helling 2019).2 Anotheradvantage of focusing on the above temperature rangeis that the upper atmospheres of such brown dwarfs areanticipated to share some resemblance with the loweratmosphere of Earth (Yates et al. 2017).

In order to calculate τBD in our simple model, we willmake use of the analytical expression derived in Burrows& Liebert (1993). Following equation (2.58) of Burrows& Liebert (1993), the effective temperature (Teff) of thebrown dwarf is given by

Teff ≈ 1551 K

(tBD

1 Gyr

)−0.324(MBD

0.05M�

)0.827

×(

κR0.01 cm2/g

)0.088

, (2)

where tBD and MBD are the age and mass of the browndwarf, respectively, and κR denotes the Rosseland meanopacity of the brown dwarf close to its photosphere. Thevalue of κR has a complicated dependence on the wave-length, density, temperature and the chemical speciespresent (Freedman et al. 2014); the weighted mixing ra-tio for cool brown dwarfs at Teff = 250 K has been pro-vided in Figure 4 of Morley et al. (2018). Owing to thevery weak dependence on κR in (2), this factor may beset aside in our order-of-magnitude analysis.

After rearranging the preceding expression and solvingfor τBD,3 we have

τBD ≈ 3.8× 105 Gyr

(MBD

M�

)2.55

, (3)

The mass range for brown dwarfs lacks precise upper andlower bounds due to ambiguities stemming from theirdefinition. Based on the orthodox definition entailingdeuterium burning, the range 0.01 .MBD/M� . 0.06 isoften employed (Caballero 2018). However, this ignoresthe existence of overmassive brown dwarfs (Forbes &Loeb 2019) as well as sub-brown dwarfs such as WISEJ085510.83-071442.5 whose mass is < 0.01M� (Luhman2014). In this paper, we will employ the range 0.005 <MBD/M� < 0.07, which is mostly coincident with theconventional limits delineated earlier.

Next, we consider the volume VBD encompassed by theatmospheric habitable zone. As the atmosphere consti-

2 It must however be noted that water vapor can exist in theatmospheres of bodies with higher effective temperatures; for in-stance, water vapor lines have been detected in Arcturus (Rydeet al. 2002), whose surface temperature is 4290 K.

3 We opt to normalize MBD in units of M� because we will dealwith stellar habitable zones subsequently. However, in the latersections, we will revert to the conventional normalization factorof Jupiter’s mass (MJ ).

3

tutes a spherical shell, we have

VBD ≈ 4πR2BDHBD, (4)

with RBD being the radius of the brown dwarf that isgiven by equation (2.37) of Burrows & Liebert (1993):

RBD ≈ 3.5R⊕

(MBD

M�

)−1/3

, (5)

where we have opted to normalize RBD in units of R⊕because we will subsequently study Earth-sized planetsin the habitable zones of stars. Next, we must assess thecharacteristic value of HBD, the vertical layer in whichhabitable conditions exist. Given that extremophiles onEarth have been documented at temperatures rangingfrom Tmin ∼ 258 K to Tmax ∼ 395 K (Clarke 2014;McKay 2014), these limits can be utilized to determineHBD. Of course, in using these limits, we are work-ing with the implicit assumption that extraterrestriallife has the same thermal limits as organisms on Earth.While this premise is indubitably geocentric, it is plau-sible that the thermal range for extraterrestrial life, asset by generic physicochemical constraints, is not far re-moved from Tmin and Tmax (Bains et al. 2015; Vladilo& Hassanali 2018). A subtle, but important, distinc-tion is worth pointing out here: Tmin and Tmax refer tothe local atmospheric temperature in the habitable zone,whereas the range for Teff introduced at the beginningof Section 2.1 corresponds to the blackbody temperatureof the brown dwarf as it cools.

From inspecting the pressure-temperature diagrampresented in Fig. 6 of Morley et al. (2014) for cool Ydwarfs, it is found that the pressure varies by a factorof ∼ 10 across this temperature range; in quantitativeterms, the pressure in this region is typically ∼ 0.1-1bar. This pressure range is compatible with the pres-sure range tolerated by Earth-based organisms (McKay2014; DasSarma & DasSarma 2018). For an isothermalatmosphere in hydrostatic equilibrium, the above resultfor the pressure implies that HBD would be a few timesthe scale height of the brown dwarf. As the scale heightis inversely proportional to the surface gravitational ac-celeration g, we specify

HBD ∼ 2.4× 10−2 km

(MBD

M�

)−5/3

(6)

after making use of equation (2.51) from Burrows &Liebert (1993) with the normalization obtained fromSection 2.1 of Yates et al. (2017).

The last term that we need to tackle is dNBD, whichcan be simplified to (dN/dMBD) dMBD. In other words,we wish to determine ξBD ≡ dN/dMBD, but this remainsdifficult to estimate, as the abundance of brown dwarfsis not rigorously constrained. We will make use of theinitial mass function described in Thies & Kroupa (2007)and Kroupa et al. (2013) that is in reasonable agreement

with numerical simulations and empirical observations(Thies et al. 2015); it has the form

ξBD ≈ 9× 10−2 C(MBD

M�

)−0.3

, (7)

where C represents a normalization constant that dropsout of our subsequent analysis. However, for late-typebrown dwarfs (spectral classes T and Y), the mass func-tion is better described by a power-law exponent of −0.6(Kirkpatrick et al. 2019); using this scaling changes oursubsequent results by a factor of ∼ 3. Therefore, aftersubstituting all of the preceding relations into (1), wecan determine ΓBD.

Before moving ahead, a comment on putative organ-isms in atmospheric habitable zones of brown dwarfs isin order. On the one hand, due to the gravitationalsettling, the microbes drift downward at the terminalvelocity. On the other, convective updraft will act totransport microbes in the opposite direction. Yates et al.(2017) developed numerical and analytical models tak-ing into account both of these factors and found thatorganisms up to ∼ 10 times larger than typical Earth-based microbes (with sizes of ∼ 1 µm) might exist asstable populations in the atmospheric habitable zonesof cool Y dwarfs. Yates et al. (2017) found that strongconvection (i.e., higher velocities) favors the evolutionof larger organisms, and vice-versa.

2.2. Circumstellar habitable zones aroundmain-sequence stars

The next category of systems we consider are Earth-sized worlds in the habitable zones (HZs) of main-sequence stars (Dole 1964; Kasting et al. 1993; Koppa-rapu et al. 2013; Ramirez 2018). As mentioned earlier,we are interested only in potential habitability, owing towhich we consider all worlds that are Earth-sized situ-ated in the HZ, regardless of whether they actually hostliquid water or not. We will use the subscript ‘?’ to in-dicate that we are studying worlds in the HZs of stars.

The first term of interest to us is τ?, namely, the pe-riod over which the planet resides in the HZ. In theory,the upper bound is determined by the stellar lifetime,but the actual value is lower because the planet will ex-perience a greenhouse effect and possibly end up beingdesiccated (Caldeira & Kasting 1992; Goldblatt & Wat-son 2012). The duration of temporal habitability wasinvestigated by Rushby et al. (2013) for planets receiv-ing Earth-like insolation. The results were subsequentlyparametrized as simple scaling relations by Lingam &

4

Loeb (2019a), which we will adopt herein as follows:

t? ∼ 5.5 Gyr

(M?

M�

)−2

M? > M�,

t? ∼ 5.5 Gyr

(M?

M�

)−1

0.5M� < M? < M�,

t? ∼ 4.6 Gyr

(M?

M�

)−1.25

M? < 0.5M�, (8)

where M? is the mass of the host star. Note that t? mayrepresent an upper bound on the habitability lifetime ofM-dwarf exoplanets. Such planets are subject to intensestellar winds (Lingam & Loeb 2017, 2018a; Dong et al.2017, 2018, 2019) and elevated X-ray and extreme UVradiation fluxes (Luger & Barnes 2015; Bolmont et al.2017), and could therefore be depleted of their atmo-spheres and water over sub-Gyr timescales; for a recentreview of this subject, see Lingam & Loeb (2019b).

The habitable volume for an “Earth-like” planet inthe HZ is found via

V? ≈ 4πR2⊕H⊕, (9)

where the habitable “height” H⊕ needs to be deter-mined. Habitability is modelled as being restricted tothe regions of the planet where the temperatures are be-tween Tmin and Tmax. This includes not only the surfaceof the planet but also its atmosphere and lithosphere. Aswe are dealing with worlds that resemble Earth in termsof their geochemical characteristics, we are free to useEarth as a proxy. First, let us consider the atmospherichabitable zone of the Earth. Based on the thermal limitsand the atmospheric temperature profile (Jacob 1999),we find that ∼ 5 km in the troposphere, ∼ 10 km nearthe stratopause, and ∼ 10 km in the thermosphere meetthis criterion. Next, if we choose an average geothermalgradient of ∼ 2.5×10−2 K/m, we find that the habitabledepth is ∼ 4 km by starting with a surface temperatureof ∼ 290 K. Thus, after adding up these contributions,we end up with H⊕ ∼ 29 km.

The last quantity to determine is ξ? ≡ dN/dM? sincedN? = η⊕ξ? dM?. The additional factor η⊕ accounts forthe number of Earth-sized planets in the HZs of theirhost stars. The estimated value of η⊕ is uncertain bynearly an order of magnitude (Kaltenegger 2017). Wewill treat η⊕ as being roughly independent of M?, andwill adopt the conservative choice of η⊕ ∼ 0.1, eventhough recent studies indicate that η⊕ . 0.35 for K-and G-type stars (Zink & Hansen 2019). To maintainconsistency with the previous section, we adopt the sameinitial mass function (Thies & Kroupa 2007; Thies et al.2015), thus yielding

ξ? ≈ 3.2× 10−2 C(M?

M�

)−1.3

0.07M� < M? < 0.5M�,

ξ? ≈ 1.6× 10−2 C(M?

M�

)−2.3

M? > 0.5M�. (10)

With all of the factors assembled, we are now in aposition to calculate Γ? after drawing upon (1). Thelower bound for the stellar mass is chosen to be 0.07M�whereas the upper bound is 2M�; increasing the upperbound to infinity alters our results by < 20%.

2.3. Ratio of habitable volumes

The chief quantity of interest is the following ratio:

∆V ≡ΓBD

Γ?. (11)

Clearly, ∆V � 1 implies that most of the habitable vol-ume is concentrated in the atmospheres of brown dwarfsrelative to that encompassed by Earth-like planets in thehabitable zone and vice-versa. After simplifying ΓBD

and Γ? by utilizing the results from the preceding sec-tions, we arrive at ∆V ∼ 52. Therefore, this result thatimplies the maximum potentially habitable volume isin the atmospheres of brown dwarfs and not Earth-likeplanets in HZs of stars.

At this stage, we note that life around stars is notconstrained to exist only in the HZ. As the examples ofEuropa and Enceladus illustrate, life may also exist insubsurface oceans underneath icy envelopes. Estimatingthe ratio of the four-dimensional volumes for Earth-likeand subsurface worlds is a much more challenging en-deavor, owing to which we shall not address this ques-tion in detail. However, the following points should beborne in mind with regards to this matter.

• It has been estimated that the maximal numberof worlds with subsurface oceans that are largerthan Europa is & 103 times higher than Earth-likeplanets in the HZ (Lingam & Loeb 2019c).

• The volume of oceans in such subsurface oceanworlds might be a few times higher with respect toEarth depending on their water inventory; Europa,in particular, may possess ∼ 2.5 times more waterthan Earth provided that its ocean depth is ∼ 100km (Chyba 2000; Lunine 2017).

• Even in the absence of tidal heating, subsurfaceocean worlds of the type specified above could re-tain their oceans over Gyr timescales via radio-genic heating (Spohn & Schubert 2003).

Based on these considerations, it is conceivable thatthe total habitable volume spanned by subsurface oceanworlds is a couple of orders of magnitude higher thanEarth-like planets in the HZ. In this event, the total vol-ume spanned by this category would be comparable tothat encompassed by the atmospheres of brown dwarfs.

3. THE PROSPECTS FOR LIFE IN BROWNDWARF ATMOSPHERES

Hitherto, we have restricted ourselves to the poten-tially habitable volume covered by brown dwarf atmo-spheres. However, a well-known fact is that habitability

5

requires more than just the appropriate thermal lim-its and the existence of liquid water (Lammer et al.2009; Shields et al. 2016; Lingam & Loeb 2019b). Wewill thus explore certain aspects pertaining to habitabil-ity and putative life in the atmospheres of cool browndwarfs. In most instances, as we possess no knowledgewhatsoever of the putative organisms that can inhabitbrown dwarf atmospheres, faute de mieux, we draw uponanalogs from Earth and adapt them accordingly.

3.1. Biomass in atmospheric habitable zones

Next, we estimate the upper bound on the possiblebiomass in brown dwarf atmospheres. Naturally, thelower bound is trivially zero in the event that life isruled out altogether in aerial settings. While we are notaware of any organisms that complete their entire lifecycles in Earth’s atmosphere, comparatively few studieshave been undertaken and additional research is nec-essary (Smith 2013). Although permanent ecosystemsin Earth’s stratosphere are “unlikely” due to the harshenvironment (Smith et al. 2010), the conditions in itstroposphere are less extreme. Moreover, for reasons doc-umented a few paragraphs below, the atmospheric hab-itable zones of cool brown dwarfs might be more clementthan the Earth’s stratosphere in some respects. Finally,the possibility of these habitable zones being seeded bymicrobes from elsewhere ought not be discounted alto-gether as per some hypotheses (Wainwright et al. 2006).

In order to estimate the biomass in brown dwarf at-mosphere, we will utilize data derived from the Earth’satmosphere. In doing so, we are implicitly invoking astrong version of the Copernican Principle, whose valid-ity is by no means confirmed. A more realistic treatmentwould attempt to derive the biomass based on nutrient(and energy) availability, but carrying out this calcula-tion requires an in-depth knowledge of the limiting nu-trients and their coupled biogeochemical cycles. Hence,we opt to use empirical results from the Earth, owingto which the ensuing findings should be viewed with ap-propriate caution.

We denote the characteristic number density of mi-crobes in the atmosphere by ρm, thus yielding a totalbiomass (Mbio) of

Mbio ∼ 4πρmR2BDHBD. (12)

As per Table 1 of Frohlich-Nowoisky et al. (2016), theaverage global mass density of microbes is ρm ∼ 10−10

kg/m3. However, in localized regions, the concentrationof airborne microbes can be as high as ∼ 1011 m−3 (Am-ato et al. 2007), which translates to a mass density ofρm ∼ 10−4 kg/m3 after presuming that the mass of asingle microbe is ∼ 10−15 kg.

Both of the above estimates pertain to modern-dayEarth. Instead, let us turn our attention to ArcheanEarth, which may have possessed a thick haze cover.From Figure 1 of Arney et al. (2016), we find that the

global aerosol density ranges from ∼ 105 m−3 in thetroposphere to ∼ 1010 m−3 in the thermosphere; theparticle radius of the aerosols is a few times smallerthan 1 µm. On modern Earth, around 20-30% of allaerosols with sizes > 0.5 µm have been documented tohost microbes (Bowers et al. 2012; Frohlich-Nowoiskyet al. 2016). Hence, if we adopt a similar fraction of 25%for Archean Earth, we find that the biomass density ischaracterized by a range of ρm ∼ 2.5× 10−11 kg/m3 toρm ∼ 2.5 × 10−6 kg/m3. Note that we have neglectedfactors other than the aerosol density to estimate theatmospheric biomass density.

In actuality, other factors such as the access to nutri-ents and thermodynamic disequilibrium as well as envi-ronmental parameters will play a major role in determin-ing atmospheric biomass density. However, it is worthappreciating that certain microbes on Earth can survivein the stratosphere despite the presence of considerabledesiccation, cold temperatures, nutrient deficiency, andhigh ultraviolet and ionizing radiation fluxes (Harrisonet al. 2013; DasSarma & DasSarma 2018). In contrast,the atmospheric habitable zones might have access towater and nutrients (the latter is described later), re-ceive lower doses of radiation due to clouds and hazes(Morley et al. 2014; Helling 2019), and (by definition) becharacterized by temperatures within the thermal limitsof Earth-based life.

In the clouds of Venus, theoretical calculations by Li-maye et al. (2018) suggest that the maximum biomassloading for aerosols with sizes > 2 µm is ρm ∼ 1.4×10−5

kg/m3. As Jupiter is the closest analog that we pos-sess for a brown dwarf in our Solar system, it is worthexamining the Jovian atmosphere in more detail. Thenephelometer carried by the Galileo spacecraft detectedparticles throughout its descent from ∼ 0.5 to ∼ 12 bars.The mean particle radius ranged between ∼ 0.5-4.0 µmand the observed number density spanned ∼ 1.4 × 105

to 3.4×107 m−3 (West et al. 2004); the maximum num-ber density was detected at a pressure of 0.75-1.35 bar,which is close to the Earth’s surface pressure. Saturnalso has aerosols in its atmosphere, whose characteris-tics are broadly similar to Jupiter (West et al. 2009).

Finally, we note that Titan’s atmosphere comprises athick organic haze that is primarily produced by photo-chemistry (Cable et al. 2012; Horst 2017). Constraintsderived from the Descent Imager/Spectral Radiometeronboard the Cassini-Huygens probe apparently indicatethat the number density of aerosols at an altitude of 80km is ∼ 5×106 m−3 with a possible mean particle size of∼ 0.7-2 µm (Tomasko et al. 2008). Photochemical andelectrical discharge experiments conducted to simulateTitan’s atmosphere have yielded higher aerosol densitiesof 4.6× 109 to 3× 1012 m−3, albeit at smaller sizes (< 1µm), as seen from Table 1 of Horst & Tolbert (2013).

Determining the aerosol density of brown dwarfs is noteasy as it depends upon the chemical species present aswell as the spectral type and altitude. Under the as-

6

10 20 30 40 50 60 7010-6

10-4

0.01

1

100

104

MBD/MJ

MaximumbiomassinEarthunits

Figure 1. The maximum biomass that may exist in cool

brown dwarf atmospheres (normalized by Earth’s biomass)

as a function of the brown dwarf mass (MBD) in units of

Jupiter’s mass (MJ). The dashed, dotted-dashed and dot-

ted curves correspond to the biomass obtained by using a

biomass density of ρmax, ρ and ρmin, respectively. The solid

horizontal line corresponds to the case where the atmospheric

biomass is equal to that of Earth.

sumptions of hydrostatic equilibrium and complete con-densation, the upper bound on the column mass densityof the particles (σc) is (Marley & Robinson 2015):

σc ≈ f(mc

m

)(Pcg

), (13)

where mc is the mass of the condensate particle, m rep-resents the mean molecular weight of the atmosphere,f denotes the mixing ratio for the condensible species,and Pc is the pressure at which condensation occurs.However, note that this formula only yields the columndensity and not the number density.

If we assign the same values of the maximum and min-imum biomass density to brown dwarfs by drawing uponthe prior examples from our Solar system, we find thatthe lower bound is ρmin ∼ 2.5 × 10−11 kg/m3 whereasthe upper bound is ρmax ∼ 10−4 kg/m3. By comput-ing the geometric mean of these two quantities, we endup with ρ ∼ 5 × 10−8 kg/m3. We are now in a posi-tion to calculate (12) for these choices of the biomassdensity. However, it is more instructive to estimate thetotal biomass normalized to that of the Earth, i.e., weintroduce the ratio

δbio ≡Mbio

Mbio,⊕. (14)

In determining Mbio,⊕, we observe that the total amountof biogenic carbon on our planet is 5.5 × 1014 kg (Bar-On et al. 2018). In order to convert the carbon content

to biomass, we multiply the former by a factor of ∼ 2(Ma et al. 2018) because the overwhelming majority ofbiomass on Earth occurs as land plants (Bar-On et al.2018). Therefore, the total biomass on Earth is esti-mated to be Mbio,⊕ ∼ 1.1× 1015 kg.

We have plotted (14) as a function of the brown dwarfmass in Figure 1 for the biomass densities delineatedpreviously. It is seen that the maximum biomass thatcan be sustained is a monotonically decreasing functionof MBD; this occurs because the radius and height ofthe habitable layer both decrease with the mass. It isalso evident from Figure 1 that the maximum biomassis actually higher than that of the Earth in all instanceswhen the optimistic value for the biomass density (ρmax)is employed. However, if the lower limit for the biomassdensity (ρmin) is utilized, the resultant biomass is ordersof magnitude smaller than Earth’s biosphere, albeit stillhigh viewed in absolute terms. The case with the meanbiomass density (ρ) is notable as straddles both regimes,i.e., δbio > 1 and δbio < 1, with the transition betweenthem occurring at MBD ≈ 6.5MJ .

3.2. Aerosols and the origin of life

There is a great deal that remains unknown aboutthe origin of life and the habitats in which it could havebeen actualized (Schulze-Makuch & Irwin 2008; Luisi2016; Kitadai & Maruyama 2018). Many of the com-mon geochemical environments posited for the originof life on Earth, such as hydrothermal vents, geother-mal fields and beaches, are clearly unavailable in at-mospheres. It has been hypothesized since the 1950sthat aerosols represent viable sites for the initiation ofabiogenesis (Goldacre 1958). This subject has witnessedsome notable developments in recent times, as reviewedin Donaldson et al. (2004) and Griffith et al. (2012).

We will briefly highlight some of the salient advantagesstemming from aerosols in a qualitative fashion, beforetackling a couple of quantitative aspects later.

• Observations and laboratory studies have revealedthat aerosols with inverted micelle structures ex-ist near water-air interfaces on Earth. Aerosols ofthis type comprise liquid water, minerals and smallorganic molecules enclosed within an organic filmmade up of fatty acids (Donaldson et al. 2004).

• These structures are akin to vesicles, although thelipid bilayers that constitute the boundary in vesi-cles possess greater functionality. Vesicles havebeen posited to play an important role in the ori-gin of protocells as they not only offer a naturalcompartmentalization scheme, but also permit thereplication of biopolymers in their interiors (Chen& Walde 2010; Blain & Szostak 2014).

• As the aerosols traverse the atmosphere, theyshould experience fluctuations in the ambient hu-midity. In brown dwarf atmospheres that may

7

have patchy water clouds (Morley et al. 2014),moving in and out of these clouds would inducea similar effect. The advantage of heterogeneity inrelative humidity is that the onset of hydration-dehydration cycles is feasible (Tuck 2002). Theimportance of such wet-dry cycles is well estab-lished, as they enable the selection, concentrationand polymerization of prebiotic monomers (Damer& Deamer 2015; Becker et al. 2018).

• The synthesis of biopolymers is obviously an im-portant step toward the origin of life. It hasbeen shown that the formation of peptide bonds,which is disfavored on thermodynamic and kineticgrounds in aqueous environments, can take placeat air-water interfaces such as those found in at-mospheric aerosols (Griffith & Vaida 2012).

As noted above, aerosols that possess the inverted mi-celle structure might contribute to the origin of proto-cells. However, it is a well-known fact that cells divide.In order to mimic protocells in this regard, it is there-fore necessary for the aerosols to be capable of fission.It was pointed out in Donaldson et al. (2001) that thefission of pure aerosol particles is not possible on ther-modynamic grounds because the free energy is alreadyat a minimum, but splitting can occur in the case ofaerosols with organic films. If the surface tension of theparent and two daughter aerosols is denoted by γP , γ1

and γ2, respectively, Donaldson et al. (2001) found thatfission could occur provided that

γ1 + γ2ζ2 < γP

(1 + ζ3

)2/3, (15)

where ζ = r2/r1 is the ratio of the radii of the daughteraerosols and has a range of (0, 1). Therefore, symmetricdivision with ζ ≈ 1 requires γ1 + γ2 < 1.59γP , whichis somewhat unrealistic as it calls for sizable changesin the surface tension of daughter particles. However,asymmetric division with 0.1 < ζ < 1 is more feasible.

Lastly, we turn our attention to the issue of spatial andtemporal scales for abiogenesis. Conventionally, abio-genesis has been envisioned as the successful outcome ofa very large number of random “trials”. Both intuitivelyand mathematically, it can be shown that the probabil-ity of abiogenesis is linearly proportional to the numberof trials (NT ) conducted (de Duve 2005). However, thisansatz is based on the premise that the mean qualityof each trial is comparable across different worlds. Aswe will henceforth compare trials in aerosol reactors inEarth’s atmosphere with those in atmospheric habitablezones of brown dwarfs, we make the assumption that themean quality per trial is similar in both instances. In as-sessing NT , it must be noted that we have to account forboth spatial and temporal factors. We begin by evalu-ating the spatial aspect. In an influential mathematicalmodel, Dyson (1999) proposed that∼ 1010 droplets werenecessary for the emergence of a low-entropy state from

10 20 30 40 50 60 705×1033

6×1033

7×1033

8×1033

9×1033

1×1034

MBD/MJ

Numberoftrials

Figure 2. The number of abiogenesis trials in aerosols that

can occur in the atmospheric habitable zones of brown dwarfs

as a function of the brown dwarf mass (MBD) expressed in

units of Jupiter’s mass (MJ). We have assumed the fidu-

cial values of na ∼ 107 m−3 and τC ∼ 100 s in (16). In

comparison, the number of trials that took place in Earth’s

atmospheric aerosols at the time life originated is ∼ 1031.

a disordered population of prebiotic monomers, analo-gous to the ferromagnetic phase transition.

The number density of aerosols (na) in cool browndwarfs will clearly be spatially variable, but we adopta fiducial value of ∼ 107 m−3 as it is compatible withobservations of Jupiter’s atmosphere (see Section 3.1).Therefore, the number of spatial trials possible in the at-mospheric habitable zone is VBDna/1010. Next, we turnto the temporal element. Dyson (1999) suggested thatthe duration of the transition to an ordered state couldbe collectively viewed as constituting a single “cycle”.Let us denote the time required for an individual cycleby τC . Thus, the number of temporal trials over the en-tire habitability interval is given by τBD/τC . Using theestimate for the number of temporal trials from Tuck(2002) in conjunction with the timescale for abiogenesison Earth based on the current fossil and phylogeneticevidence (Dodd et al. 2017; Betts et al. 2018), we ob-tain a fiducial value of ∼ 100 s for the Earth after divingthe latter by the former. By combining the prior factorstogether, we obtain

NT ∼ 4× 1033

(MBD

MJ

)0.22 ( na107 m−3

)( τC102 s

)−1

.(16)

In contrast, it has been estimated that NT . 1031 ifone selects the time at which the first evidence for lifeon Earth appears in the geological record (Tuck 2002).Therefore, we find that the maximum number of tri-als available to brown dwarfs over their entire habitableperiod is possibly ∼ 3 orders of magnitude higher with

8

respect to Earth’s atmosphere at the time of abiogenesis,as seen from Figure 2 where NT is plotted as a functionof MBD for the choices na ∼ 107 m−3 and τC ∼ 100 s.

3.3. Nutrient and energy availability

Apart from the necessity of liquid water, both energyand nutrients are known to be essential preconditions forthe origin and sustenance of biospheres. Modeling theenergy and nutrient inventories is a complex endeavor,owing to which we shall focus on a few select aspects.

3.3.1. Bioessential elements

The chief bioessential elements are CHNOPS. We willnot explicitly deal with carbon, hydrogen and oxygen asthey are clearly accessible in the form of atmosphericmethane and water, respectively. The issue of sul-fur metabolism has been investigated in the context ofVenus (Schulze-Makuch et al. 2004; Limaye et al. 2018).In particular, it has been suggested that the analogs ofAcidithiobacillus ferrooxidans could exist on Venus. Inanaerobic conditions, A. ferrooxidans uses Fe3+ as anelectron acceptor and oxidizes elemental sulfur to gen-erate products such as sulfuric acid.

With regards to sulfur, one of the key points worthbearing in mind is that virtually all of the sulfur inven-tory in brown dwarf atmospheres is anticipated to existin the form of hydrogen sulfide (Visscher et al. 2006). Inaddition, sulfide clouds comprising ZnS, MnS and Na2Sare expected to exist. Therefore, as the most commonversion of sulfur exists in the form of sulfide, it wouldmake sense for putative metabolic pathways to utilizethe sulfide anion. One of the best known metabolicpathways involving H2S as the reactant is anoxygenicphotosynthesis, with Chlorobiaceae (green sulfur bacte-ria) representing a classic example. A schematic repre-sentation of this pathway is

CO2 + 2H2S→ CH2O + H2O + 2S. (17)

However, two difficulties encountered by this reactionare the availability of CO2 and photons of suitable en-ergy; note, however, that sufficient CO2 may exist inthe lower atmosphere (Bilger et al. 2013) and we ad-dress photon availability in Section 3.3.2.

Sulfur oxidizing bacteria represent another interest-ing candidate (Nordstrom & Southam 1997), but theytypically require electron acceptors such as oxygen andnitrate (NO−3 ), neither of which are abundant in coolbrown dwarf atmospheres. On the other hand, metaloxides present in the atmosphere might instead play therole of reactants and undergo reduction. Before movingon, we note that trace quantities of FeS could exist, buthigh abundances are unlikely as iron is expected to con-dense first to yield Fe clouds (Visscher et al. 2006, 2010).Note that iron-sulfur compounds are highly important inlife-as-we-know-it, because they constitute the basis of

iron-sulfur clusters found in many key proteins (Bandy-opadhyay et al. 2008), and may have placed a vital rolein the origin of life (Wachtershauser 1990).

Next, we turn our attention to nitrogen. It is well-known that ammonia (NH3) is a prominent componentof the atmospheres of cool brown dwarfs. The existenceof ammonia obviates the necessity for biological nitrogenfixation (biosynthesis of ammonia) by diazotrophs. Theavailability of ammonia, or ammonium (NH+

4 ) in someinstances, is also vital from the standpoint of Earth’s ni-trogen cycle because aerobic ammonia oxidizing bacte-ria convert ammonia to nitrite through a compound re-action involving hydroxylamine (Kowalchuk & Stephen2001) illustrated below:

2H+ + NH3 + 2e− + O2 → NH2OH + H2O

NH2OH + H2O→ NO−2 + 5H+ + 4e−. (18)

Subsequently, nitrite is converted to nitrate by nitriteoxidizing bacteria; in turn, nitrate is acquired by biolog-ical organisms as it constitutes a vital nutrient. In theabsence of oxygen, anammox (anaerobic ammonium ox-idation) bacteria can oxidize ammonia to yield N2. Onaccount of the presence of atmospheric ammonia, mi-crobial metabolic pathways along the above lines maybe feasible, but they will necessitate the availability ofsuitable oxidants such as oxygen or nitrite.

The last major bioessential element that needs to beaddressed is phosphorus (Westheimer 1987). In the caseof both modern and Proterozoic Earth, dissolved phos-phorus (in the form of phosphates) is often regardedas the ultimate limiting nutrient over long timescales(Tyrrell 1999; Sarmiento & Gruber 2006; Laakso &Schrag 2018). The same conclusion is expected to holdtrue for certain classes of exoplanets such as oceanworlds (Lingam & Loeb 2019d). One of the chief issueswith phosphorus availability is that most phosphorus onEarth exists as mineral phosphates, which are character-ized by very low solubilities; the best known example inthis category is fluorapatite (Ca5(PO4)3F). At a pH of 7and a temperature of 307 K, the solubility of fluorapatitein pure water is ∼ 3× 10−3 g/L (McCann 1968).

We will now pivot to phosphorus content in cool browndwarfs. At the effective temperatures considered herein,under the assumption of equilibrium chemistry, most ofthe phosphorus should exist in the form of tetraphospho-rus hexaoxide (P4O6) as per theoretical calculations (Fe-gley & Lodders 1994). Of more interest to us is the com-pound ammonium dihydrogen phosphate (NH4H2PO4).Its condensation temperature (Tc) is estimated fromequation (52) of Visscher et al. (2006) as:

Tc ≈104 K

30− 0.2 (11 logPT + 15[X/H]), (19)

where PT is the pressure (in units of bar) and [X/H]quantifies the metallicity. As an example, for a brown

9

dwarf with solar metallicity at an altitude wherethe pressure is ∼ 1 bar, the condensation tempera-ture is found to be Tc ∼ 333 K. It is therefore ap-parent that cool brown dwarfs might host clouds ofNH4H2PO4. Based on the observed L band spectra ofthe Y dwarf WISE J085510.83-071442.5, Morley et al.(2018) suggested that the detected obscuration of near-infrared flux could be explained through the presence ofNH4H2PO4 clouds at a pressure of ∼ 10 bar.

As the effective temperature of WISE J085510.83-071442.5 (∼ 250 K) is within the thermal range consid-ered herein, it appears likely that other brown dwarfswith this effective temperature also possess NH4H2PO4.The reason we have underscored the existence ofNH4H2PO4 has to do with the issue of phosphoruslimitation delineated above. The first, and perhaps themost important point, to appreciate is that NH4H2PO4

is very soluble in water, thereby yielding the dihydrogenphosphate anion that may be utilized by putative or-ganisms. At room temperature (∼ 300 K), the solubilityof ammonium dihydrogen phosphate is ∼ 4 × 102 g/L(Lide 2007), about 105 times higher than fluorapatite.

Second, we note that NH4H2PO4 has been widely em-ployed in laboratory experiments of prebiotic chemistry.A few salient examples are listed below.

• In the 1960s and 1970s, several studies establishedthat the phosphorylation of nucleosides to yieldnucleotides and their oligomers (precursors of nu-cleic acids) was facilitated through the additionof ammonium dihydrogen phosphate and heat-ing the mixtures (Ponnamperuma & Mack 1965;Lohrmann & Orgel 1971; Horowitz & Hubbard1974; Oro & Stephen-Sherwood 1976).

• Glycerol phosphates, which are important precur-sors of complex lipids that comprise cell mem-branes, can be synthesized by heating a mixtureof glycerol and ammonium dihydrogen phosphate(Epps et al. 1979; Deamer & Oro 1980). It shouldalso be noted that the synthesis of phosphate am-phiphiles (analogous to phospholipids in cell mem-branes) has been facilitated by using NH4H2PO4

(Powner & Sutherland 2011).

• In view of the ubiquity of polyphosphates (e.g.,adenosine triphosphate) in biology, by draw-ing upon prior analyses, Keefe & Miller (1995)pointed out the fact that such compounds could(in)directly be generated by heating NH4H2PO4

at fairly moderate temperatures of ∼ 333-373 K.

Although a number of laboratory experiments employedNH4H2PO4, they either did not endeavor or were unableto identify plausible sources for this compound (Keefe& Miller 1995). The latter aspect is not surprisingbecause ammonium phosphates (including dihydrogenphosphate), despite their intrinsic advantages, were un-likely to have been available on Hadean Earth (Pasek

et al. 2017). In contrast, as we have seen previously, coolbrown dwarfs may have ammonium dihydrogen phos-phate clouds at pressures of ∼ 1-10 bar that constitutepotentially viable sources of this compound.

Finally, we note that a number of other bioessentialelements are predicted to exist in brown dwarf atmo-spheres such as manganese and iron. Hence, nutrientlimitation vis-a-vis these elements might not necessarilypose a serious issue. However, we are not aware of anyempirical or theoretical constraints on the abundance ofmolybdenum (Mo) or tungsten (W) in brown dwarf at-mospheres.4 Molybdenum, in particular, is an essentialcomponent of many enzymes in prokaryotes and eukary-otes (Schwarz et al. 2009), with the most notable beingthe nitrogenases that reduce nitrogen to yield ammonia;certain prokaryotes have evolved to use tungsten in lieuof molybdenum (Hille 2002). Hence, if Mo (or W) is abioessential element insofar as life-as-we-know-it is con-cerned, it is important to gauge the abundance of thiselement in brown dwarf atmospheres.

3.3.2. Electromagnetic energy

There are a number of energy sources that are accessi-ble for prebiotic synthesis as well as putative biospheresin brown dwarf atmospheres.5 Some examples includecosmic rays, radioactivity, lightning and chemical energy(Deamer & Weber 2010; Lingam & Loeb 2019c).

Instead of quantifying the fluxes for all these sources,we will focus only on electromagnetic radiation. OnEarth, solar radiation constitutes the primary sourceof energy. It is therefore not surprising that photosyn-thesis is responsible for the majority of carbon fixationand biomass on Earth (Bar-On et al. 2018). The emer-gence of photosynthesis was one of the major evolution-ary transitions in our planet’s history (Knoll 2015). Onaccount of these reasons, we opt to analyze the avail-ability of photosynthetically active radiation (PAR) inthe atmospheres of cool brown dwarfs.

Before proceeding further, a comment regarding thenature of photosynthesis is in order. The most pro-ductive version of photosynthesis on Earth is oxygenicphotosynthesis, but its feasibility is probably loweredfor primarily anoxic atmospheres. Instead, in the H2-dominated atmospheres of brown dwarfs, it is plausiblethat “hydrogenic photosynthesis” is a viable pathway(Bains et al. 2014). The net reaction is expressible as

CH4 + H2O→ CH2O + 2H2, (20)

4 The same issue also applies to the Venusian atmosphere, aspointed out by Lingam & Loeb (2018b).

5 In connection with prebiotic synthesis, owing to qualitativesimilarities between the atmospheres of cool brown dwarfs andJupiter, it is conceivable that some prebiotic compounds listed inSection 1 for the latter may also be synthesized in the former.

10

which is more transparent when separated into the twoconstituent half-reactions given by

CH4 + H2O→ CH2O + 4H+ + 4e−

4H+ + 4e− → 2H2. (21)

One of the major benefits underlying hydrogenic pho-tosythesis is that both methane and water are abun-dant in cool brown dwarf atmospheres (Lodders & Feg-ley 2002; Cushing et al. 2011; Zahnle & Marley 2014).There are a number of possible advantages stemmingfrom hydrogenic photosynthesis that were reviewed inBains et al. (2014). Two of the most pertinent ones are:(a) the energy required for synthesizing a given quantityof biomass is ∼ 5-10 times lower than oxygenic photo-synthesis, and (b) the longest wavelength suitable for hy-drogenic photosynthesis is 1.5 µm, whereas conventionaloxygenic photosynthesis ostensibly requires photons atwavelengths of . 750 nm (Nurnberg et al. 2018).

We begin our analysis by considering free-floatingbrown dwarfs that do not receive any radiation fromexternal sources. As the atmospheric altitudes underconsideration are smaller than RBD, to leading orderwe can model the maximum available PAR flux by theblackbody flux emitted in the appropriate wavelengthrange. The photon flux (ΦBD) for the blackbody is

ΦBD =

∫ λmax

λmin

2c

λ4

[exp

(hc

λkBTeff

)− 1

]−1

dλ, (22)

where we work with Teff ≈ 300 K since we are ana-lyzing brown dwarfs with effective temperatures in therange ∼ 250-350 K. The maximum (λmax) and mini-mum (λmin) photon wavelengths corresponding to PARare difficult to determine for other worlds, and will beaddressed shortly hereafter. For now, we adopt the lim-its λmin ≈ 0.35 µm and λmax ≈ 1.1 µm as these wave-lengths are compatible with those used by anoxygenicphotoautotrophs on Earth (Kiang et al. 2007a; Ritchieet al. 2018). In our subsequent analysis, we will holdλmin fixed as this is consistent with photosynthesis be-ing inhibited by UV radiation (Hollosy 2002; Caldwellet al. 2007; Castenholz & Garcia-Pichel 2012).

After substituting the above values into (22), we endup with ΦBD ≈ 1.2×106 m−2 s−1. In contrast, the mini-mum photon flux required for photosynthetic organismson Earth is Φc ≈ 1.2× 1016 m−2 s−1 (Raven et al. 2000;Wolstencroft & Raven 2002). This lower bound been ex-plained through physicochemical constraints imposed byH+ leakage, charge recombination, and protein turnover.Hence, as per the conventional PAR limits, it is impos-sible for photosynthesis to function in the atmospheresof free-floating brown dwarfs.

Next, it is instructive to consider the total emittedphoton flux (Φtot) given by

Φtot ≈ 4.8c

(kBTeff

hc

)3

. (23)

Evaluating Φtot for Teff ≈ 300 K, we obtain Φtot ≈1.3 × 1022 m−2 s−1. Interestingly, this value is ∼ 6.5times higher than the total solar photon flux incidenton the Earth. Therefore, insofar as total photon flux isconcerned, the lower atmospheres of brown dwarfs couldreceive a higher photon flux relative to the Earth. Thechief difference, however, is that most of the photons areemitted at wavelengths of ∼ 10 µm.

We will now examine the maximum wavelength thatis necessary to facilitate oxygenic photosynthesis. Wol-stencroft & Raven (2002) (see also Hill & Rich 1983)suggested that oxygenic photosynthesis could functionat longer wavelengths by harnessing a higher numberof long-wavelength photons to carry out the fixation ofCO2, which is consequently accompanied by the genera-tion of O2. In the event that multiple (> 2) photons arerequired per electron, it is important to recognize thatcertain components of the hypothesized photosyntheticapparatus (e.g., water-oxidation complex) will probablydiffer from those found on Earth in several key respects,such as having a multiple Z-scheme (Kiang et al. 2007b).As the number of photons required increases, the evo-lution of a more intricate mechanism with its attendantcatalytic macromolecules becomes necessary. It is im-possible to ascertain whether the evolution of this appa-ratus is practically feasible; however, insofar as physicsis concerned, there are no manifestly evident constraintsthat appear to rule it out.

The following theoretical relationship has been pro-posed for multi-photon oxygenic photosynthesis (Wols-tencroft & Raven 2002; Kiang et al. 2007b):

χ ≈ 2

(λmax

0.7µm

), (24)

where χ denotes the number of photons required perelectron transfer in oxygenic photosynthesis; the nor-malization is specified based on oxygenic photoau-totrophs on Earth. In this event, the minimum photonflux required becomes (χ/2) Φc. By imposing the con-straint that ΦBD should equal this value and invoking(22), we obtain λmax ≈ 2.7 µm and χ ≈ 7.8. In otherwords, in order for oxygenic photosynthesis to occur, aunique photosystem based on eight photons per electronwould presumably be necessary; for χ = 8 (the given χshould be an integer), we find λmax = 2.8 µm.

At this stage, some crucial caveats should be men-tioned. First, our analysis presupposed that all photonsin the range λmin < λ < λmax will reach the atmosphericlayer where the biota are present. Second, the minimumphoton flux represents an ideal limit as it assumes that100% absorption by the photosystem(s). Third, if eightphotons are required per electron transfer, the absorp-tion of an equal number of photons at higher energiescould cause overheating and disrupt the photosyntheticapparatus. Finally, we note that thermodynamics it-self places strict constraints on the efficiency at whichthe radiation can be utilized. Suppose that the ambient

11

temperature in the atmospheric layer is Ta. The Carnotefficiency (ηC) serves as an upper bound in most (butnot all) instances.6 The standard expression for ηC is

ηC = 1− TaTeff

. (25)

Considering Ta ≈ 280K and Teff ≈ 310 K, we see thatthe Carnot efficiency is only around 10%. In fact, theabove formula predicts that extracting work becomesimpossible when Ta > Teff . A more accurate treatmentnecessitates calculating the exergy in the PAR range(Delgado-Bonal 2017; Scharf 2019), which goes beyondthe scope of this paper. The Carnot efficiency constraintbecomes relatively unimportant when it comes to coolbrown dwarfs that are companions of stars.

On account of the above limitations, it is very plau-sible that the above values of λmax and χ derived con-stitute lower bounds. On the other hand, we note thatthe photosynthetic machinery in other worlds might pos-sess a higher efficiency and functionality with respect toEarth-based organisms as a result of having evolved inlow-light conditions. In addition, our derivation ignoredthe possibility of PAR being derived from a host star(if one exists) or even high-energy astrophysical objectssuch as active galactic nuclei (Lingam et al. 2019).

It is worth examining the PAR accessible from thehost star in more detail. By modelling the star as ablackbody, it is possible to estimate the critical orbitalradius (ac) at which the PAR flux becomes equal to Φcintroduced previously. After simplifying the resultantexpression, we obtain

ac ∼ 316 AU

(L?L�

)1/2(T?T�

)−1/2√I(T?)

I(T�), (26)

where L? and T? are the bolometric luminosity and tem-perature of the host star, whereas I is given by

I(T?) ≈∫ `2(T?)

`1(T?)

x′2 dx′

exp (x′)− 1, (27)

where we have introduced `1(T?) ≈ 3.32 (T?/T�)−1

and

`2(T?) ≈ 7.12 (T?/T�)−1

. In deriving (26), we haveopted for the conservative PAR range of 0.35-0.75 µm.In actuality, as noted earlier, the maximal wavelengthcould increase if multi-photon schemes are viable. Theonly difference in this case is that I(T?) in (26) needs tobe replaced by the function K(T?) defined as

K(T?) ≈2

χ

∫ `2(T?)

2`1(T?)/χ

x′2 dx′

exp (x′)− 1, (28)

where it should be recalled that χ denotes the numberof photons used per electron transfer.

6 Note that the efficiency of a Carnot engine presented in (25)is not always correct (Curzon & Ahlborn 1975).

χ = 2

χ = 3

χ = 4

0.5 1.0 1.5 2.0 2.5

1

10

100

1000

M/M⊙

Criticalorbitalradius(AU)

Figure 3. The maximum orbital radius at which photosyn-

thesis might be feasible as a function of the stellar mass (M?)

in units of solar mass (M�). The three curves correspond

to the number of photons (χ) utilized per electron transfer,

where χ = 2 for conventional oxygenic photosynthesis.

In Figure 3, we have plotted the critical orbital radiusac as a function of the stellar mass (M?) after employ-ing empirical mass-radius and mass-temperature scal-ings described in Lingam & Loeb (2019e) as well as thenumber of photons (χ) involved in photosynthesis. Frominspecting Figure 3, we see that the number of photonsutilized per electron transfer does not alter our resultssignificantly for stars with M? > M�. However, whenwe consider M? ∼ 0.1M�, we find ac ≈ 1.9 AU for con-ventional photosynthesis (χ = 2), whereas ac ≈ 4.9 AUfor χ = 4. Hence, the evolution of multi-photon schemescould enable an increase of the width of the photosyn-thesis “zone” around low-mass stars.7

Our estimate for ac in (26) is merely a loose upperbound because we have not taken the opacity of thebrown dwarf atmosphere in the visible (or near-IR) intoaccount. In reality, (26) must be multiplied by the fac-tor exp (−τV ), where τV quantifies the optical depth ofthe brown dwarf atmosphere (until the atmospheric hab-itable zone is reached) in the appropriate wavelengthrange. Hence, it is possible that ac might be smallerby orders of magnitude but this is hard to determine apriori since the optical depth is regulated by pressureand temperature, the abundances of various chemicalspecies, and the wavelength (Marley & Robinson 2015).In the event that most of the PAR from the star is ob-scured by the upper atmosphere when it reaches the

7 We have not explicitly investigated the role of stellar flares inpowering photosynthesis, as their contribution is probably mini-mal for the majority of stars (Lingam & Loeb 2019e).

12

habitable zone, the outcome for photosynthesis is ren-dered analogous to free-floating brown dwarfs that wehave previously analyzed in this Section.

4. DETECTING LIFE IN BROWN DWARFATMOSPHERES

We will briefly examine brown dwarf statistics, dis-cuss potential biosignatures in the atmospheres of browndwarfs, and the prospects for detecting them.

4.1. Brown dwarf statistics

Until this stage, we have not explicitly addressed thequestion of whether the brown dwarfs under consider-ation are situated around a host star, in binaries, orfree-floating. We will briefly explore this issue, as ithas implications for the search strategies. An inspec-tion of (2) reveals that achieving an effective tempera-ture of Teff ∼ 300 K within the current age of the Uni-verse (∼ 1010 yr) is feasible only for brown dwarfs withMBD . 20MJ . Hence, broadly speaking, late-type Tdwarfs and Y dwarfs are of primary interest to us.

In the 2000s and 2010s, surveys of brown dwarf bina-ries have established the following properties: (a) theiroccurrence rate (. 20%) is lower than the stellar popula-tion, (b) they are mostly found in tightly bound orbits atseparations of a few AU, and (c) their mass ratio distri-bution has a sharp peak near unity and declines rapidlythereafter (Close et al. 2003; Burgasser et al. 2006; Kraus& Hillenbrand 2012). In a recent study, Fontanive et al.(2018) surveyed late-type T dwarfs and Y dwarfs andfound that the binary fraction was ∼ 5.5%, with a peakin separation at 2.9 AU and a power-law exponent of 6.1for the mass-ratio distribution.

Next, we consider brown dwarf companions to stars.It has been suspected since the 1980s (Campbell et al.1988) that there is a paucity of brown dwarfs within afew AU of solar-type stars (Campbell et al. 1988), whichcame to be known as the brown dwarf “desert” (Marcy& Butler 2000). This was confirmed by a number ofsubsequent studies of FGKM stars (Evans et al. 2012;Ma & Ge 2014; Reggiani et al. 2016), which found thatthe fraction of brown dwarfs at distances of ∼ 1-100 AUwas merely a few percent (Metchev & Hillenbrand 2009;Sahlmann et al. 2011; Dieterich et al. 2012; Cheethamet al. 2015), whereas giant planets were relatively moreabundant by a factor of . 10 (Grether & Lineweaver2006; Lafreniere et al. 2007).

A statistical analysis of stars from various spectralclasses at different ages concluded that the probabil-ity distribution function (P) of substellar companionsobeyed P ∝ M−0.65a−0.85, where M and a were themass and orbital radius of the substellar object; it wasfound that the fraction of stars hosting 5-70MJ objectsat 10-100 AU was ∼ 1-3% at 68% confidence (Brandtet al. 2014). Although close-in brown dwarfs are rela-tively scarce, it does not imply that they are completelyabsent. For instance, nine brown dwarf companions

to solar-type stars have been identified at distances of∼ 0.5-4 AU based on data collected by the SOPHIEspectrograph (Wilson et al. 2016).

4.2. Potential biosignatures

The biosignatures produced would be directly depen-dent on the putative organisms in question. For exam-ple, two of the most well-known biosignatures on Earthare molecular oxygen (O2) and the “red edge” of vege-tation, both of which are consequences of oxygenic pho-tosynthesis (Schwieterman et al. 2018).

First, we consider the presence of dead or decompos-ing organisms. In this scenario, in lieu of the organ-isms themselves their constituent biomolecules wouldplay a prominent role. In the case of many Earth-based biomolecules, it is well-known that their peak ab-sorption lies in the UV and visible regions. A reviewof many of these biomolecules can be found in Limayeet al. (2018). Nucleic acids and proteins have peak ab-sorbances at wavelengths of 260 nm and 280 nm, respec-tively. Iron-sulfur proteins, which play vital roles in re-dox reactions, are characterized by maximum absorptionat wavelengths of 280-450 nm. A number of biologicalpigments such as chlorophylls, carotenoids and pterinsare strong absorbers at wavelengths < 500 nm. Thus,viewed collectively, we might expect to see an increasein the reflectance (i.e., decrease in absorption) spectraat wavelengths & 500 nm. As this behavior has beendocumented for the Venusian atmosphere, it has led tosuggestions that Venus’ clouds might be harboring mi-crobes (Limaye et al. 2018).

Next, let us turn our attention to live organisms. Inprinciple, a number of microbial metabolisms are fea-sible as outlined in Section 3.3.1, but some variant ofphotosynthesis comes across as a natural candidate dueto its ubiquity and importance on Earth. Along theexpected lines, the spectral red edge roughly coincideswith λmax on Earth. If we posit that a similar situationholds true for the photoautotrophic organisms in coolbrown dwarf atmospheres, various possibilities open updepending on whether the brown dwarf is free-floatingor bound as well as the nature of the photosyntheticpathway.

We begin by tackling photosynthesis on free-floatingbrown dwarfs. As explained in Section 3.3.2, λmax for(oxygenic) photosynthesis is ∼ 2.8 µm. If the minimumflux required for hydrogenic photosynthesis is compara-ble to its oxygenic counterpart, we can solve for λmax,with 0.7 µm in the denominator of (24) replaced by 1.5µm, after postulating that the biophysical process re-quires two photons of 1.5 µm. Hence, we end up withλmax ≈ 2.6 µm and χ ≈ 3.5; in other words, hydrogenicphotosynthesis entailing 3-4 photons might be feasible.For χ = 3 and χ = 4, the corresponding values of λmax

are approximately 2.25 µm and 3.0 µm, respectively.Therefore, for free-floating brown dwarfs, the manifes-tation of a “red edge” close to the outer boundary of the

13

near-infrared (near-IR), i.e., at wavelengths of ∼ 2.3-3.0µm, comes across as being plausible.

Next, we turn our attention to brown dwarf compan-ions around stars. As long as the criterion a < ac issatisfied, enough photons for photosynthesis (either hy-drogenic, anoxygenic or oxygenic) should be accessibleto photoautotrophs. Hence, if the exact analog of oxy-genic photosynthesis exists, we would expect to see aspectral edge at ∼ 700 nm. Instead, it is more plausiblethat the spectral edge will be manifested in the near-IRat wavelengths of . 1.5 µm, as both anoxygenic and hy-drogenic photoautotrophs can utilize such wavelengths(Kiang et al. 2007b; Bains et al. 2014).

Apart from the photosynthetic spectral edge, the de-tection of biosignature gases produced by oxygenic pho-tosynthesis is challenging. The oxygen thus producedwould react quickly with reduced gases (that are abun-dant) unless it is generated in large quantities. Ammo-nia is a potential byproduct of hydrogenic photosynthe-sis (Bains et al. 2014), but it will be very challenging todistinguish between biotic and abiotic NH3, given thatthe latter is plentiful in cool brown dwarfs. However,an interesting avenue for possibly identifying hydrogenicphotosynthesis stems from noting that methane is de-pleted in accordance with (20).

Hence, if there is a mismatch between the abundanceof methane inferred through observations and that de-termined by theory using only abiotic sources and sinks,it might be indicative of biological activity. Cooler at-mospheres exhibit stronger signs of disequilibrium and,in principle, the diagnosis of chemical disequilibrium canbe undertaken via the analysis of second eclipse spectra(Line & Yung 2013; Krissansen-Totton et al. 2016); seealso Krissansen-Totton et al. (2018). The spectra of theY dwarf WISE J085510.83-071442.5 are compatible withan under-abundance of methane (Morley et al. 2018),but we emphasize that this discrepancy (if it exists) isexplainable via abiotic mechanisms.

Needless to say, even the detection of such features isnot indicative of life because it may instead arise fromfalse positives. For instance, several abiotic materialssuch as dust, salts and polymers have been argued toexplain the reduction in Venusian albedo observed at. 500 nm, with sulfur dioxide and iron chloride beingtwo notable candidates (Zasova et al. 1981). Likewise,minerals such as cinnabar (HgS) display sharp spectraledges that are reminiscent of the red edge of vegetation,albeit not at the same location (Seager et al. 2005).Along similar lines, it is conceivable that some of thespectral edges elucidated above overlap with those pro-duced by abiotic substances.

Second, as the atmospheres of brown dwarfs compriselayers of clouds, their existence will hinder measure-ments. However, in the event that the cloud cover ispatchy - a feature that has been confirmed for somebrown dwarfs (Radigan et al. 2012) - time-resolved spec-tra might permit the identification of spectral features.

Microscopic organisms may also experience horizontaland vertical transport due to atmospheric circulation(Showman & Kaspi 2013; Zhang & Showman 2014; Apaiet al. 2017), and could therefore be transported to re-gions with lower opacity, thereby presumably renderingtheir spectral features more discernible.

4.3. Detectability of biosignatures

As remarked previously, brown dwarfs can either ex-ist on their own (i.e., free-floating) or as stellar com-panions. Observing free-floating brown dwarfs is advan-tageous from the standpoint of not having to concernourselves with resolving their spectra by subtracting thecontribution from the host star.

On the other hand, the emission peak of brown dwarfsat Teff ∼ 300 K is at ∼ 10 µm, which typically en-tails observations undertaken in the mid-IR. For exam-ple, detailed spectra of WISE J085510.83071442.5 (withTeff ∼ 250 K) have been obtained in the L and M bands,corresponding to wavelength ranges of 3.4-4.14 µm and4.5-5.1 µm, respectively. Yet, as outlined in Section 4.2,several interesting biosignatures are expected to mani-fest at visible and near-IR wavelengths, where the browndwarf is many orders of magnitude fainter.

For cool Y dwarfs at distances of . 20 pc, the firstupper limits on the abundances of gases such as CH4,H2O, NH3, H2S and CO2 were obtained recently us-ing the Wide Field Camera 3 instrument on the HubbleSpace Telescope (Zalesky et al. 2019). The Near In-fraRed Spectrograph (NIRSpec) on the upcoming JamesWebb Space Telescope (JWST) operates over a wave-length range of 0.6-5 µm.8 Numerical simulations un-dertaken by Zalesky et al. (2019) suggest that a signal-to-noise ratio (SNR) of ∼ 200 in the J-band is achievablewith only ∼ 15 minutes of integration time for objectsat distances of ∼ 10 pc, implying that it represents apowerful tool for characterizing the atmospheres of Ydwarfs. Hence, searching for atmospheric biosignaturesof free-floating cool brown dwarfs is viable with JWST.A quantitative estimate of the yield from JWST is de-scribed toward the end of the section.

Next, directing our attention to brown dwarf compan-ions, it was noted in Section 4.1 that there exists a browndwarf desert relative to giant planets. At this stage, werecall that cool brown dwarfs and giant planets shareseveral similarities, although there also exist apprecia-ble differences (e.g., surface gravity). Hence, our subse-quent discussion is also applicable to giant planets withmasses M > MJ that may possess aerial biospheres. Wewill adopt the scaling relations reviewed in Winn (2010)and Fujii et al. (2018) henceforth.

Even though (5) is more accurate for brown dwarfs,we will utilize it to calculate the radius R of both browndwarfs and giant planets a few times more massive than

8 https://jwst.nasa.gov/nirspec.html

14

Jupiter; we refer to them collectively as substellar ob-jects. First, we note that the transit depth scales as R2,implying that the transit depth of the substellar objectis ∼ 102 times higher with respect to an Earth-sizedplanet. Next, for transmission spectroscopy, the SNRmanifests the scaling:

SNR ∝ RH∆t1/2, (29)

where ∆t is the integration time and H is the scaleheight of the substellar object. We have opted to holdthe stellar properties constant as well as the instrumentspecifications. Now, if we wish to determine the inte-gration time for a fixed SNR, we see that

∆t ∝( a

RH

)2

. (30)

Using the fact that the scale height is H ∼ kBTa/(mg),the above expression reduces to

∆t

∆t⊕∼ 2.85× 10−7

(M

MJ

)4(Ta

250 K

)−2

, (31)

after using equation (2.51) of Burrows & Liebert (1993)for g. Note that ∆t⊕ denotes the integration time re-quired for an Earth-like planet. Hence, if we chooseTa ∼ 250 K and M ∼ 10MJ , we see that the integra-tion time required to achieve a particular SNR drops bythree orders of magnitude for this substellar object.

Next, if one wishes to study thermal emission (i.e.,emission spectrum) from the substellar object, the SNRwill scale as

SNR ∝ R2∆t1/2, (32)

where the other parameters are held constant. Now, asbefore, if we consider a fixed value of SNR, we end upwith ∆t ∝ R−4, which simplifies to

∆t

∆t⊕∼ 6.3× 10−7

(M

MJ

)4/3

. (33)

Therefore, upon selecting M ∼ 10MJ , from the aboveformula we see that the integration time relative to anEarth-sized planet decreases by five orders of magnitude.

Lastly, let us suppose that we are interested in di-rect imaging of a cool substellar object via reflectedlight. The contrast ratio scales as R2, implying thatit increases by a factor of ∼ 102 in comparison to anEarth-like planet. The SNR ought to exhibit the samescaling as (32) apart from an extra factor of a−2, im-plying that the desired integration time is given by (33)when a is held fixed. Hence, the corresponding integra-tion time would be lowered by a factor of ∼ 105 withrespect to an Earth-sized planet, when we consider asubstellar object with M ∼ 10MJ .

Thus, the basic conclusion to be drawn herein is thatthe integration times required are considerably reduced,

implying that achieving a high SNR is orders of mag-nitude more feasible when dealing with substellar ob-jects with masses M & 10MJ relative to characterizingEarth-like exoplanets. We will now quantify the yieldof cool brown dwarf companions to stars, whose atmo-spheres are analyzable by JWST.

Out to a distance of d? ∼ 100 pc from Earth, thereare ∼ 5× 105 stars, of which ∼ 50% are M-dwarfs withM? . 0.2M� (Kroupa et al. 2013). Considering an or-bital radius of ∼ 5-30 AU (with a geometric mean of∼ 12 AU), it can be assumed that the fraction of starswith brown dwarfs is on order of 1% (Evans et al. 2012;Dieterich et al. 2012). If we further specialize to objectswith M < 20MJ , the yield must be further lowered bya factor of ∼ 3 based on the substellar IMF specified inKirkpatrick et al. (2019). Thus, by combining all thesefactors, we find that ∼ 103 cool brown dwarfs might besuitable for biosignature characterization by JWST.

We begin with the case of characterizing substellarobjects via transmission spectroscopy. We make useof equation (4) of Fujii et al. (2018) derived for JWSTand work with H ∼ 20 km (Showman & Kaspi 2013),9

R ∼ RJ (RJ is Jupiter’s radius) and d? ∼ 100 pc. Theresultant SNR for JWST is found to be

SNR ∼ 3.3

(∆t

1 hr

)1/2

, (34)

implying that a moderately high SNR is achievable forcool substellar objects even with hours of integrationtime. Next, we turn our attention to detecting the emis-sion spectrum from these objects. We make use of equa-tion (7) of Fujii et al. (2018) for the above choice ofparameters, and obtain

SNR ∼ 2.5ε

(∆t

1 hr

)1/2

, (35)

where ε embodies the relative depth of spectral fea-tures. Hence, if we choose an integration time of a fewhours, a reasonable SNR is attainable. A potential issuewith detecting emission from brown dwarf companions isthat the starlight needs to be separated from the browndwarf, owing to which it is probably easier to study free-floating brown dwarfs via this avenue.

The last mode of observation entails reflected lightfrom substellar objects. However, owing to the 1/a2 de-pendence and the relative paucity of brown dwarfs atdistances of a few AU, this method is disfavored com-pared to the previous two methods described herein. Fora ∼ 10 AU and Bond albedo of ∼ 0.2, the required con-trast ratio to differentiate the substellar object from thestar is ∼ 10−10, which is challenging with state-of-the-art coronographs and starshades (Fujii et al. 2018).

9 It must be noted that the scale height is not constant for browndwarfs because it is dependent upon the pressure and composition(Marley & Robinson 2015).

15

5. CONCLUSION

The search for extraterrestrial life outside our Solarsystem is expected to play a major role in the near-future. Currently, virtually all theoretical and obser-vational studies are geared toward finding atmosphericbiosignatures of rocky planets in the habitable zones oftheir host stars. However, despite a few studies in thecontext of our Solar system, the potential for life in at-mospheric habitable zones (i.e., aerial biospheres) hasmostly gone unappreciated, with perhaps the only note-worthy exception being Yates et al. (2017).

In this paper, we have therefore investigated the at-mospheric habitability of cool brown dwarfs, as well assub-brown dwarfs and giant planets, at an effective tem-perature of ∼ 250-300 K. In Section 2, we began byestimating that the maximum habitable volume encom-passed by cool brown dwarf atmospheres is conceivablytwo orders of magnitude higher than the volume asso-ciated with Earth-like planets in the habitable zones oftheir host stars. The reasons for the higher habitablevolume are the greater spatial volume and temporal du-ration that collectively offset the fact that stars are morenumerous than brown dwarfs.

As there are many facets of putative extraterrestrialaerial biospheres that have not been investigated hith-erto, we explored some of the key aspects in Section3. By drawing upon data for Earth’s current aerialbiosphere in conjunction with empirical constraints onother Solar system objects, we found that the biomassencapsulated in the atmospheric habitable zones of coolbrown dwarfs might surpass the Earth’s biomass underoptimal conditions. Next, we highlighted the signifi-cance of aerosols as potential prebiotic reactors in fa-cilitating the origin of life, and thereby showed that thenumber of abiogenesis “trials” in brown dwarf aerosolspossibly exceeds that of the Earth at the time of life’sappearance around 4 Ga by a factor of ∼ 100-1000.10

We surveyed the bioessential elements accessible toputative organisms in the atmosphere. We directed mostof our attention toward phosphorus, as it constitutes thelimiting nutrient on Earth. We highlighted the forma-tion of ammonium dihydrogen phosphate and how itcould serve as a ready source of soluble phosphorus aswell as yield a number of vital prebiotic compounds.Subsequently, we explored the prospects for photosyn-thesis on free-floating brown dwarfs. Despite the gen-eral paucity of photons, we hypothesized that photosyn-thesis could function via a multi-photon scheme with amaximum wavelength of ∼ 2.3-3.0 µm, i.e., close to the

near-IR outer boundary. In contrast, for brown dwarfsthat are stellar companions, photon availability is notexpected to be a major limiting factor in most instancesand the spectral edge would occur at either visible ornear-IR wavelengths depending on the stellar spectrum.

In Section 4, we presented a brief survey of browndwarf statistics and assessed the spectral biosignaturesthat can result from the presence of life. We proposedthat the analog of the red edge of vegetation might oc-cur, albeit at wavelengths in the near-IR; the exact loca-tion of the spectral edge is dependent on the metabolicpathway. Another possibility is that chemical disequi-librium could result from the depletion or generationof certain gases (e.g., methane) that is potentially de-tectable. In the case of cool substellar objects aroundstars, we demonstrated that the required integrationtime to achieve a high SNR is orders of magnitudesmaller with respect to Earth-sized planets at roughlythe same distance and effective temperature. We findthat ∼ 103 cool brown dwarfs may be investigated forbiosignatures by JWST at distances . 100 pc, with anintegration time of O(1) hr yielding a SNR of ∼ 5.

Thus, viewed collectively, there is arguably a strongcase to be made for seeking atmospheric biosignaturesin cool brown dwarfs and sub-brown dwarfs.11 A ma-jor advantage with pursuing this line of enquiry is thateven the non-detection of life will still provide us withan in-depth understanding of planetary atmospheres be-cause such objects exhibit stellar composition but areotherwise akin to giant planets in their atmosphericphysics and chemistry. Apart from observational sur-veys, laboratory experiments are needed to properlygauge whether life could exist in conditions mimick-ing these cool atmospheres and what types of biosig-natures would be most prominent. Finally, laboratoryexperiments and observations must be supplanted withtheoretical and numerical models that assist in makingtestable predictions and interpreting empirical results.

We thank the reviewer for the very comprehensive andinsightful report that helped improve the quality of thepaper. We also thank Freeman Dyson and Corey Powellfor valuable comments regarding the paper. This workwas supported in part by the Breakthrough Prize Foun-dation, Harvard University’s Faculty of Arts and Sci-ences, and the Institute for Theory and Computation(ITC) at Harvard University.

10 Our hypothesis concerning the origination of life within theatmospheres of cool substellar atmospheres ignores the possibilityof these worlds being seeded by way of interstellar panspermia,despite the fact that recent calculations suggest that it might befeasible (Belbruno et al. 2012; Lingam 2016; Ginsburg et al. 2018).

11 If life is detected on these worlds someday, perhaps they willmerit the sobriquet “green dwarfs”. The word “green” is partic-ularly apropos if the existence of chlorophyll-type photosyntheticpigments is revealed, as the green color of vegetation on Earth isa direct consequence of chlorophylls.

16

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