# Venusian Habitable Climate Scenarios: Modeling Venus through time and
applications to slowly rotating Venus-Like Exoplanets

###### Abstract

One popular view of Venus’ climate history describes a world that has spent
much of its life with surface liquid water, plate tectonics, and a stable
temperate climate. Part of the basis for this optimistic scenario is the high
deuterium to hydrogen ratio from the Pioneer Venus mission that was
interpreted to imply Venus had a shallow ocean’s worth of water throughout
much of its history. Another view is that Venus had a long lived ($\sim$ 100
million year) primordial magma ocean with a CO2 and steam atmosphere. Venus’
long lived steam atmosphere would sufficient time to dissociate most of the
water vapor, allow significant hydrogen escape and oxidize the magma ocean. A
third scenario is that Venus had surface water and habitable conditions early
in its history for a short period of time ($<1$Gyr), but that a moist/runaway
greenhouse took effect because of a gradually warming sun, leaving the planet
desiccated ever since. Using a general circulation model we demonstrate the
viability of the first scenario using the few observational constraints
available. We further speculate that Large Igneous Provinces and the global
resurfacing 100s of millions of years ago played key roles in ending the
clement period in its history and presenting the Venus we see today. The
results have implications for what astronomers term “the habitable zone,” and
if Venus-like exoplanets exist with clement conditions akin to modern Earth we
propose to place them in what we term the “optimistic Venus zone.”

JGR: Planets

NASA Goddard Institute for Space Studies, 2880 Broadway, New York, NY, 10025,
USA GSFC Sellers Exoplanet Environments Collaboration Theoretical
Astrophysics, Department of Physics and Astronomy, Uppsala University,
Uppsala, SE-75120, Sweden

M.J<EMAIL_ADDRESS>
Venus could have had habitable conditions for nearly 3 billion years.

Surface liquid water is required for any habitable scenario.

Solar insolation through time is not a crucial factor if a carbonate-silicate
cycle is in action.

## Plain Language Summary

We have little data on our neighbor Venus to help us understand its climate
history. Yet Earth and Venus are sister worlds: they initially formed close to
one another, and have nearly the same mass and radius. Despite the differences
in their current atmospheres and surface temperatures, they likely have
similar bulk compositions, making comparison between them extremely valuable
for illuminating their distinct climate histories. We analyze our present data
on Venus alongside knowledge about Earth’s climate history to make a number of
exciting claims. Evaluating several snapshots in time over the past 4+ billion
years, we show that Venus could have sustained liquid water and moderate
temperatures for most of this period. Cloud feedbacks from a slowly rotating
world with surface liquid water reservoirs were the keys to keeping the planet
clement. Contrast this with its current surface temperature of 450 degrees and
an atmosphere dominated by Carbon Dioxide and Nitrogen. Our results
demonstrate that it was not the gradual warming of the sun over the eons that
contributed to Venus’ present hot-house state. Rather we speculate that large
igneous provinces and the global resurfacing 100s of millions of years ago
played key roles in ending the clement period in its history.

##  1 Introduction

The case for prolonged habitability of Venus has been made by a number of
authors in recent years [<]e.g.,¿[]GrinspoonBullock2007,Way2016. If so, then
if habitability was widespread and persisted over geological timescales (e.g.
109 years), it is possible that organisms were capable of filling a large
variety of environmental niches as occurred on Earth via evolutionary
processes. This has led to speculation about possible remaining life-friendly
niches [<]e.g.,¿[]MorowitzSagan1967,Grinspoon1997,Cockell1999,Schulze-
Makuch2002,Schulze-Makuch2004,Dartnell2015,Limaye2017. These possibilities
alone provide sufficient motivation to understand whether early Venus was
habitable or not. However, Venus is also interesting from the perspective of
the large number of rocky exoplanets discovered to date inside the traditional
inner edge of the habitable zone but far enough from their host stars to
maintain an atmosphere [Kane . (2019)]. If some of these planets in close
proximity to their host stars have long periods of habitability it may
overturn traditional notions of the habitable zone [<]e.g.,¿[]Kopparapu2013
and influence target selection for characterization of the atmospheres of
these planets. Hence, what appears to be a modern Venus-like world close to
its parent star might host surface liquid water. We refer to such habitable
worlds as residing in an “Optimistic Venus Zone,” a subset of the planets in
the “Venus Zone” described in Kane2014.

There are three primary requirements for the early habitability of Venus. The
first is that temperatures were low enough that liquid water was capable of
condensing on the surface of Venus. The second is that Venus had a sufficient
inventory of water to create the conditions believed necessary for the rise of
life on terrestrial worlds [<]e.g.,¿[]Brack2010. Third, volatile cycling and a
geologically active surface must exist to regulate the climate as possibly
supported by the work of Noack2012. For surficial water reservoirs most
research has focused on the poorly constrained measurement of the D/H ratio of
$\sim$ 150$\pm$30 (times Earth’s value) by the Pioneer Venus Mission
(https://nssdc.gsfc.nasa.gov/planetary/pioneer_venus.html) [Donahue . (1982),
Donahue . (1997)]. A number of other ground based measurements have been made
[De Bergh . (1991), Marcq (2006), Bjoraker . (1992), Fedorova . (2008)] that
have generally revealed values of D/H greater than 100. For a recent review of
D/H measurements of Venus see Section 6.3 of Marcq2017. We review the work of
authors who have considered Venus’ climate evolution below.

The first simple gray radiative transfer calculations of Venus’ climate
history suggested an early runaway greenhouse effect
[<]e.g.,¿[]Ingersoll1969,Rasool1970. Any water would have remained in vapor
form throughout Venus’ early history until its loss to space via
photodissociation processes [Goody Walker (1972), Walker (1975)]. CO2 would
continue to be outgassed and would accumulate (offset by atmospheric loss over
the last 4.5Gyr) to the values we see today.

Pollack1971 used a similar model but with non-gray radiative transfer. This
was probably the first work to show that Venus could have hosted liquid water
on its surface in its early history, but they also demonstrated that a runaway
greenhouse was possible as well. The difference between the two outcomes was
_mostly_ due to to the specified cloud cover and water vapor content of the
atmosphere. With 50% cloud fraction the result was usually a runaway
greenhouse. For an N2 dominated atmosphere with a modern Earth water vapor
profile, 100% cloud cover, and a 30% less luminous sun (akin to $\sim$4.5Ga)
he achieved surface temperatures $\sim$300K. These purely radiative models did
not include convection, which would limit the water vapor mixing ratio by
precipitation and limit the lapse rate of temperature to the moist adiabatic
value [Rampino Caldeira (1994)] and lessen the strength of the greenhouse
effect to some degree. Likewise, they did not include an atmospheric general
circulation that would determine the cloud fraction self-consistently.
DonahuePollack1983 replicated some aspects of the Pollack1971 work and again
showed that clement conditions were possible in Venus’ early history.

Kasting1984 used a 1-D radiative-convective model to demonstrate that a planet
with an insolation S0X=1.45 times that of modern Earth (1973 W m2 $\sim$
3.8Gya at Venus’ orbit) would have temperatures $\sim$100∘ C for a 2 bar N2
H2O atmosphere. This atmosphere contained a wet stratosphere, and thus
photodissociation of H2O and subsequent escape of hydrogen
[<]e.g.,¿[]GoodyWalker1972 could explain the long-term loss of Venus’
primordial ocean and its present dry conditions. In a subsequent study
Kasting1988 claimed that clouds would cool the atmosphere enough to keep it in
a moist greenhouse state, rather than the runaway state. At the same time the
Kasting1988 surface temperature for his “Early Venus” (Figure 7) cloud-free
model was $\sim$500K and hovered just at the margin of the moist vs. runaway
greenhouse states. Thus it was clear early on that maintaining liquid water on
ancient Venus required high fractional cloud cover, but whether physical
mechanisms exist to produce these cloudy conditions was not addressed. These
early Venus habitability scenarios would have taken place within the first
billion years of the planet’s evolution limiting the possibility for complex
life to evolve. Venus would have subsequently entered a moist/runaway
greenhouse and left us with the state it is in today. For the curious reader
BullockGrinspoon1999 provide a nice review of the literature on the
possibility of an early habitable Venus. Note that in a previous work [Way .
(2016)] it was shown that if habitable conditions were possible in early
Venus’ history that it likely remained so and that increasing solar insolation
through time is not a deciding factor. In the present work we find the same,
and hence believe that limited (in time) early habitability models are not
supported by our simulations.

Grinspoon and Bullock, in a number of abstracts
[<]e.g.,¿[]GrinspoonBullock2003,Grinspoon2008, described a Venus climate
evolution scenario with long-lived surface habitability consistent with that
described later in this paper. This work was never published, though, so the
details of their calculations and the justifications for their conjectures are
not documented.

The first three-dimensional (3-D) General Circulation Model (GCM) exploration
of issues relevant to ancient Venus was that of Abe2011. Their study focused
on land planets (analogous to the planet Arrakis of _Dune_) with no permanent
surface water reservoirs, only limited fixed amounts of ground water. In some
scenarios they found that “in principle Venus could have been a habitable land
planet as recently as 1 billion years ago.” Their conclusions resulted from
the limited water vapor in the atmosphere caused by the modest subsurface
reservoir, its limited greenhouse effect, and the albedo of the planet.
Thereafter Leconte2013 used a 3-D GCM to look at climate scenarios for Gl581 c
and HD 85512 b. HD 85512 b orbits a K-dwarf with a _s_ ynchronous rotation
period of $\sim$ 58 days and receives approximately 1.86 times the insolation
of present day Earth. This is slightly less than present day Venus’ insolation
(1.9), but its host star is a K-dwarf rather than a G-dwarf and hence the peak
of its Planck blackbody spectrum is at longer wavelengths. They state, “if not
for its thick atmosphere, Venus’ climate would be very close to the one of HD
85512 b.” They modeled a dry planet akin to that of Abe2011 using an N2
dominated atmosphere with 376ppm CO2, and no water vapor. The model produced
large temperate regions for a number of different atmospheric pressures (See
Leconte2013 Figures 1, 2) for their land planet setup, but they also found
similar behavior in their ‘collapse’ scenario with different atmospheric water
vapor profile amounts.

A rationale for high albedo cloud cover on ancient Venus was first presented
by Yang2014. Yang2014 was a large parameter study looking at the inner edge of
the habitable zone around solar type stars using The National Center for
Atmospheric Research (NCAR) Community Atmosphere Model (CAM) with
comprehensive atmospheric physics but a thermodynamic ocean. A thermodynamic
ocean, also known as a slab or mixed-layer ocean, is typically of limited
depth ($<$100 meters) with a prescribed horizontal ocean heat transport, or no
horizontal heat transport at all. The latter implies that the temperature of
each ocean grid cell is determined solely by the atmosphere and incident
sunlight directly above it. The shallow depth reduces the time lag between
solar forcing and ocean response. See Way2017 Section 2.2.2 for details. In
essence they stepped the sidereal rotation rate of an Earth-like world from 1
to 256 sidereal days. At the same time they increased the insolation from that
of modern Earth to as much as 2.6 times the modern Earth for their most slowly
rotating world of 256 sidereal days period. One of their key conclusions was
that the slowest rotators would have had a day-night general circulation that
would generate an optically thick contiguous cloud bank in the substellar
region. This cloud deck would greatly increase the planetary albedo, keeping
the surface temperature moderate even for quite high values of insolation.
These results were later confirmed with a completely different 3-D GCM with a
fully-coupled dynamic ocean [Way . (2018)]. Yang2014 also included a
simulation with modern Venus orbital parameters, spin rate, insolation while
using modern Earth topography and land/ocean mask, but again with a
thermodynamic ocean. These studies provide a possible rationale for the cloud
cover needed to produce temperate surface conditions as first postulated by
Pollack1971. Finally ¡¿[hereafter Paper I]Way2016 took things one step further
by exploring different topographies, insolations and rotation rates to put
tighter constraints on possible habitable conditions for ancient Venus.

In this paper, we extend those parameter studies to consider a wider variety
of planets and attempt to justify our modeling assumptions in light of
possible scenarios for Venus’ evolution. Sections 2–5 review existing
observational constraints and hypotheses about the composition, thickness, and
evolution of Venus’ atmosphere and water history; its surface and interior;
and its rotation and obliquity. In Section 6 we draw upon this information to
inform a series of 3-D global climate model simulations to illustrate possible
scenarios for an early habitable Venus that transitioned to its current
inhospitable state relatively late in its history. We discuss the implications
of our results for the design of future missions to Venus and for the
potential habitability of exoplanets inside the inner edge of the traditional
“habitable zone” in Section 7. Finally, recently published complimentary work
by WellerKiefer2019 supports many of our conclusions.

##  2 Atmospheric Composition and Pressure

Assuming that Venus and Earth formed from the same parts of the protoplanetary
disk and thus with similar compositions [<]e.g.,¿[]Raymond2004, we are guided
by the history of Earth whose early atmosphere was likely CO2-rich and
possibly cool, but not frozen due to the faint young sun for the late Hadean
and early Archean (4.2-4.0 Ga)
[<]e.g.,¿[]Owen1979,Kasting1993a,ZahnleSleep2002,Valley2002,Zahnle2006,ZahnleSchaeferFegley2010,Kunze2014,CatlingKasting2017,Krissansen-
Totton2018,Mello2019,Mojzsis2019. The picture of a “Cool Early Earth” promoted
in the early 2000s by Valley2002,Valley2005 and with more recent zircon data
[Valley . (2014)] may also be applicable to Venus’ early evolutionary history
if Venus survived its magma ocean (MO) phase with some liquid surface water.
Chassefiere2012 was probably the first to make the comparison of Hadean Earth
with Venus, calling his hypothesis a “cool early Venus.”

The prospects for a cool early Earth and Venus hypothesis have improved in
recent years, while the likelihood for their end in a cataclysmic “late
instability” known as the Late Heavy Bombardment (LHB) may have started to
fall out of favor. The LHB is important because its intensity was imagined so
intense that the surface temperatures of Earth would return to those before
the cool early Earth period, perhaps evaporating the oceans and/or raising
surface temperatures to values exceeding 100∘C. However, in recent years a
number of studies have started to question the timing and strength of the LHB
[<]e.g.¿[]Boehnke2016,Zellner2017,Morbidelli2018,Quarles2019,Mojzsis2019. For
a brief overview of the latest work on the LHB see Mann2018 and Voosen2020. An
“early instability” has started to gain favor in what is termed an “accretion
tail” scenario. This accretion tail scenario has recently been promoted in the
works of ¡¿[Figure 1]Morbidelli2018 and ¡¿[Figure 1]Mojzsis2019. In these
works a late or early heavy bombardment is triggered by a period of orbital
‘instability’ experienced by the giant planets as described in Ribeiro2020.
More recent work by Ribeiro2020 demonstrates that the instability has a median
timescale of 36.78-–61.5$Myr$ and is within 136$Myr$ in 75% of their cases.
Ribeiro2020 say the timing of an early instability fits in nicely with the
survival of the Patroclus–Menoetius Jupiter Trojan as a primordial binary from
the Kuiper Belt [Nesvorny . (2018)] and is a good match to other solar system
properties [<]e.g.¿[]Clement2019, but they explicitly discuss the many
limitations of their model. In the spirit of an early instability and
accretion tail scenario, Figure 4 of Mojzsis2019 suggests that Venus accreted
less than 0.01 wt% between 4.3 and 4.1Ga, about the same as Earth. As
Mojzsis2019 notes, “Results show that an abating impact flux from late
accretion is inadequate to sterilize the surface zone.” These works
demonstrate the possibility that the cool early Earth continued through the
originally proposed time period of the LHB from $\sim$ 3.8–4.1Ga. Hence we
will begin our most ancient Venus simulations with CO2 dominated atmospheres
at $\sim$4.2Ga (during the Valley2002 cool early Earth period) that evolves
over time to N2 dominated atmospheres. If the LHB were a real event we would
want to begin our simulations at 3.8Ga, rather than 4.2Ga, as others have
speculated previously [<]e.g.,¿[]RampinoCaldeira1994,Lammer2018. There are
still many unanswered questions regarding the early post-MO history of the
atmospheres of Venus, Earth and Mars. Even though we have far more information
to discern Earth’s early post-MO atmosphere it remains a complicated story yet
to be fully resolved [<]e.g.,¿[]Hirschmann2012,Nikolaou2019 and may depend on
atmospheric pressure more than previously assumed
[<]e.g.,¿[]Gaillard2014,Bower2019.

The carbonate-silicate cycle [<]e.g.,¿[]Walker1981,Stewart2019 is the key to
keeping most of Earth’s CO2 locked up in rocks for much of its history rather
than in the atmosphere as on present day Venus. The carbonate-silicate cycle
on Earth functions via subductive-type plate tectonics, the presence of a
hydrosphere and continental crust. For Earth there are several lines of
evidence to suggest these may go back to the Hadean [<]e.g.,¿[and references
therein]Mojzsis2001,Hopkins2008,Harrison2009,Korenaga2013,Harrison2017,ONeill2017,Rozel2017,Kusky2018,Korenaga2018,Maruyama2018.
Dehant2019 reviews the literature for a later beginning of plate tectonics on
Earth (Section 3.2). The requirements for how plate tectonics begins are still
not fully understood and hence remains an active area of research
[<]e.g.,¿[]Lenardic2019. If Venus had a similar early atmospheric and interior
evolution to that of Earth then early volatile cycling via some form of plate
tectonics is a viable hypothesis. In fact recent work demonstrates convection
regimes like that of plate tectonics in Venus’ recent history [Gillmann
Tackley (2014), Davaille . (2017)]. Dehant2019 reviews plate tectonic
mechanisms that may not require as much water as previously believed, which
may be relevant if Venus retained some water but less than Earth. Hence in our
scenario a carbonate silicate cycle is hypothesized for Venus after the magma
ocean phase and well before its resurfacing period to keep CO2 largely
sequestered in crustal carbonates as on modern Earth.

We are motivated to choose 1 bar atmospheres for our epochs of interest based
on geological atmospheric pressure proxies for Earth that imply an atmospheric
density of $\sim$0.25–1.1bar for most of the past 4 billion years [Nishizawa .
(2007), Goldblatt . (2009), D. Som S. Catling . (2012), S. Som . (2016), Marty
. (2013), Avice . (2018)]. However, some models of Nitrogen cycling imply that
the atmospheric pressure could have been higher in the Archean [Johnson
Goldblatt (2015), Mallik . (2018)]. Differences in N2 of factors of a few have
relatively small net impacts on climate, primarily due to small
decreases/increases in the pressure broadening of CO2 lines and partly
offsetting decreases/increases in Rayleigh scattering. Given the likely
similar geochemistry and volatile histories of Venus and Earth
[<]e.g.,¿[]RingwoodAnderson1977,Lecuyer2000,Treiman2009,Chassefiere2012,Rubie2015,
a similar evolution in their early composition and pressure histories is
plausible. Although whether they did indeed start out with similar volatile
inventories is still an open research question [<]e.g.¿[]Horner2009.

Additional work demonstrates that the present day CO2 and N2 inventories of
Earth and Venus may be similar [Donahue Pollack (1983), Goldblatt . (2009)] if
Venus is more degassed than Earth as DonahuePollack1983 speculated. Early work
by Rasool1970 and Kasting1988 estimated nearly as much CO2 was locked up in
carbonate rocks on Earth as exist in the atmosphere of Venus. Since there is
no carbonate-silicate cycle active on Venus today most of the CO2 that would
otherwise be locked up in the interior is in the atmosphere, having degassed
over the past several hundred million years (at least). As well, nearly 4
times as much N2 is found in Venus’ present day atmosphere as in Earth’s when
scaled by planetary mass. Earth’s internal N2 budget is estimated at $\sim$ 7
$\pm$ 4 times that of the atmospheric mass [Johnson Goldblatt (2015)].

However, Argon (Ar) modeling studies [Kaula (1999), O’Rourke Korenaga (2015)]
imply that Venus is less degassed than Earth, although see DonahuePollack1983
and Halliday2013 for alternative explanations for the current 40Ar abundance
difference between Earth & Venus. Watson2007 claims that 40Ar is not a
reliable indicator of degassing, although “the interpretation of their data is
controversial” according to ORourke2015. Halliday2013 also mentions that
Venus’ radiogenic 40Ar implies Venus is less degassed than Earth, but he notes
that Venus’ atmospheric non-radiogenic 36Ar has concentrations roughly two
orders of magnitude greater than Earth’s and that “this is hard to explain
unless it (Venus) is heavily degassed with a larger inventory of primordial
volatiles.” In further support of a degassed Venus Halliday2013 notes that the
similar Earth and Venus budgets of C and N support a degassed Venus while
helping to explain the high 20Ne and 36Ar budgets of Venus. One of the largest
stumbling blocks in deciding whether the 40Ar modeling inference is correct is
a proper measurement of K/U on Venus, which is presently highly unconstrained.
For this reason modeling studies generally use Earth values.
Namiki_Solomon1998 use He to confirm the 40Ar estimates, but they require a
Venus in-situ mission to make the appropriate Th and U abundance measurements
to characterize the geochemistry necessary for He production. One also needs
to consider He escape, an under-explored area of research given that it
depends not only on the ability of He to remain charged or neutral, but also
whether Venus had a past magnetic field and how present day measurements can
distinguish in situ He in the upper atmosphere from solar wind deposition.

Another limiting factor in comparing Earth and Venus is the lack of good
constraints on the bulk water content of Earth. Some papers refer only to
surface water amounts of 0.023 wt% of the Earth’s total mass, but there are
compelling arguments for possibly deep reservoirs within the lower mantle or
core [<]e.g.,¿[]Raymond2006,Schmandt2014,Genda2016,Ikoma2018. For Venus,
interior water amount estimates are mostly speculation and are restricted to
planetary formation studies [<]e.g.,¿[]Ikoma2018.

##  3 Venus’ Early Evolution & Evidence for Water

Venus & Earth likely received similar initial water inventories during their
formation histories, as has been shown in a number of works
[<]e.g.,¿[]Raymond2006. It is also tied to their composition in general, as
discussed above in Section 2.

The Pioneer Venus mass spectrometer measured a very high D/H ratio of
150$\pm$30 times that of terrestrial water [Donahue . (1997)] in a trapped
droplet of sulfuric acid. It is the only such published in-situ measurement.
Other non in-situ measurements have been made as noted in Section 1, and work
continues apace [<]e.g.¿[]Tsang2017. The original Venus D/H discovery paper by
Donahue1982 was titled “Venus was Wet: A measurement of the Ratio of Deuterium
to Hydrogen.” This was a tantalizing prospect, but limited by our knowledge of
atmospheric escape processes [Donahue Pollack (1983), Donahue Russell (1997),
Donahue (1999)] and the D/H of delivered materials over the aeons.

It is possible that the D/H ratio is not an indicator of large amounts of
water in Venus’ ancient history. Grinspoon1993 pointed out that a short
residence time for water in the present atmosphere of Venus works against the
primordial ocean hypothesis. Grinspoon1993 also noted that updated theoretical
calculations at that time that implied higher deuterium escape efficiency put
constraints on the D/H source water of 10-15. That would rule out source
material such as meteorites, comets and dust particles with high D/H ratios
[<]e.g.,¿[]Irvine2000,Charnley2009. As well, measurements of D/H and Xenon
isotopes in Comet 67P [Altwegg . (2015), Marty . (2017)] imply that Earth’s
ocean has a much lower contribution from cometary objects than previously
thought. This would also likely rule out a large cometary contribution to the
high D/H ratio measured on Venus.

A number of authors have tried to model changes in the D/H fractionation over
time [Kasting Pollack (1983), Gurwell Yung (1993), Gurwell (1995), Hartle .
(1996)] to put some constraints on when the water was lost. Given the lack of
data from Venus it is equally difficult to constrain or move these models
forward.

The possibility that the high D/H ratio implies long-lived surface water is
also limited by in-situ measurements. There is some circumstantial evidence of
past surface water from surface emissivity observations from the Galileo NIMS
instrument [Hashimoto . (2008)] and the Venus Express VIRTIS instrument
[Mueller . (2008)]. These observations may imply that the highland “tessera”
regions are mostly composed of felsic rocks, and if they are indeed granitic
they would have required surface water to form [Campbell Taylor (1983)].
Gilmore2015,Gilmore2017 find that at least one tessera region observed with
VIRTIS (Alpha Regio) appears to be more felsic than surrounding plains. This
also suggests that these older stratigraphic units [Ivanov Basilevsky (1993),
MS. Gilmore . (1997)] are granitic crustal remnants, but recent work by
Wroblewski2019 shows that parts of the Ovda Regio highland tessera are not in
fact of granitic origin. However, it is not yet possible to generalize the
work of Wroblewski2019 to the entirety of tessera.

Nikolayeva1990 and Shellnutt2019 analyzed surface rock measurements from
Venera 8. As Shellnutt2019 explains “it is possible that the Venera 8 probe
encountered a fragment of crust that resembles a terrestrial greenstone belt.”
Zolotov1997 and Johnson2000 have also demonstrated that signatures of water in
hydrous minerals may persist on the surface of Venus for long periods even
after the surface morphology has changed. This motivates an in-situ mission to
Venus to search for such materials, if they exist. Watson2007 tried to
demonstrate that the 40Ar in Earth’s atmosphere is related to the hydration of
the oceanic lithosphere consisting of relatively Ar-rich olivine and
orthopyroxene. If the results from Watson2007 are correct, (and there is
skepticism [Ballentine (2007)]), this would lead one to believe that the 40Ar
in Venus’ atmosphere today implies that water oceans could have persisted for
some time.

An outstanding unsolved and understudied problem is what happened in the epoch
of Venus’ MO as it cooled, as this may greatly affect the long-term water
inventory of the planet. The timescale of the MO crystallization could be of
order a few million years ($Myr$) as for Earth
[<]e.g.,¿[]Katyal2019,Nikolaou2019 or greater than 100 $Myr$ [Hamano . (2013),
Lebrun . (2013)]. The longevity of the MO and associated hot steam and CO2
atmosphere is vital to understanding the volatile history of Venus
[<]e.g.¿[]Salvador2017. If the MO and steam atmosphere persist too long then
much of the primordial water inventory of Venus could have been lost in its
very early history from a stronger solar wind [Chassefière (1997),
Lichtenegger . (2016)]. An attractive feature of the extended MO hypothesis is
that it naturally solves the problem of the lack of oxygen in the present day
Venusian atmosphere. This would be accomplished by sequestering the O2 left
behind by H2O dissociation in the magmatic crust and upper mantle
[<]e.g.¿[]Lebrun2013,Gillmann2009,Lichtenegger2016,Lammer2018. We discuss how
large quantities of O2 can be lost after a significant period of habitability
in bullet 6 of Section 8, in lieu of early MO losses. It is possible that the
high D/H ratio we see today [Donahue . (1982), Donahue . (1997)] is a relic of
the early MO period. If the MO cooled quickly, then there was an opportunity
to build up a surface ocean and atmosphere as is believed to have happened in
Earth’s early history. The question is whether Venus’ surface conditions as a
result of its closer proximity to the Sun would prevent the condensation of
water on its surface or not. The answer is more complicated than it may seem
since water can condense under hot high pressure multi-bar atmospheres.
MatsuiAbe1986 allow for temperatures up to 600K, while later work by Liu2004
allow temperatures approaching 720K.

Additionally the answer may reside in the planet’s rotation history, what role
clouds played, and the outgassing rates of H2O and CO2. As we will show in
Section 5 it is possible for Venus to reach a tidally locked state in less
than a few hundred $Myr$ using constant phase lag dissipation theory,
suggesting that the planet’s rotation rate could have been slow early on. As
shown in previous work [Way . (2016), Way . (2018)] as long as a planet is in
the slowly rotating regime (length of day greater than $\sim$16 Earth sidereal
days) its climate dynamics work to allow liquid water to persist on the
surface for insolations up to $\sim$ 2.6 times that of present day Earth. This
is due a large contiguous dayside cloud deck that significantly increases the
planetary albedo as discussed in Section 4.

The timing of the MO termination is critical in more than one way. If the
steam and CO2 atmosphere cooled sufficiently for MO crystallization to occur
by the time of the Late Veneer (also referred to as “Late Accretion”) then
even if Venus lost most/all of its primordial H2O through escape processes
[Gillmann . (2009), Hamano . (2013), Lichtenegger . (2016)] there may have
been a second chance to obtain a surface ocean, albeit a shallow one. Recent
work by Greenwood2018 implies that Earth may have received as much as 30% of
its H2O inventory in post-accretion impact delivery, consistent with research
that shows that the entire H2O budget cannot come from the late veneer
[Morbidelli Wood (2015)]. Halliday2013 concludes that if veneers were common
they should be proportional to planetary mass, and hence Venus would have
received a percentage of late veneer H2O similar to that of Earth. If Venus
was left dry after a long-lived magma ocean phase [Hamano . (2013)], then this
amount of H2O veneer also fits within the error bounds of Venus’ measured D/H
ratio [Donahue . (1982), Donahue . (1997)]. It should be noted that the work
of Greenwood2018 can also fit within the Ruthenium studies of
FischerGodde2017. For a contrary point of view see Gillmann2019, who claim
that most of the late veneer impactors would have been Enstatite/ordinary
chondrites which are water-poor, as opposed to water-rich carbonaceous
chondrites that would have been a mere 0-2% of the total chondrite delivery.
These contrary points of view come about because different geochemical
measurements give different answers as pointed out in a number of recent works
[<]e.g.¿[]Albarede2009,FischerGodde2017,Dauphas2017,McCubbin2019,Zahnle2020.
There are two other important caveats to consider regarding the late veneer.
First the water content depends upon the composition of the accreting bodies.
For example, if the late veneer was made up of a few large bodies then the
variations could have been greater than if it was due to a large collection of
smaller bodies. Second, the definition of the late veneer is important since
it is typically associated with the accretion of bodies after the last giant
impact. Jacobson2017 has suggested that the last giant impact on Venus could
have been much earlier than on Earth and this obviously affects the
composition of objects making up the late veneer.

At the same time, work by Gillmann2009, Morbidelli2000, and Raymond2006 shows
that Venus’ initial water inventory at formation could be as much as two
terrestrial ocean’s worth while large planetary embryos could deliver much
more within 200 $Myr$ of formation. If true, and if the magma ocean lifetime
on Venus was shorter rather than longer, then our estimates of the water
content on Venus from Pioneer Venus D/H ratios [Donahue . (1982), Donahue .
(1997)] should be more toward the higher end, $\sim$ 16% of a present day
Earth’s ocean [Donahue . (1997)]. However, it is not clear whether Venus’
primordial water content can readily be constrained by the D/H ratio
[Grinspoon (1987), Grinspoon (1993)] due to a lack of knowledge of sources and
sinks over the lifetime of the planet.

Another hypothesis [<]e.g.,¿[]RampinoCaldeira1994 states that because of its
proximity to the Sun, Venus could never condense water on its surface and
hence its surface temperature has always been 300K or higher [<]see Figure 1
in¿[]RampinoCaldeira1994 and that most of this water was lost by
photodissociation [Goody Walker (1972)]. At the same time the lack of water
prevents silicate rock weathering (on Earth this removes CO2 from the
atmosphere), hence the CO2 builds up in the atmosphere driving temperatures
ever higher due to the greenhouse effect as seen today.

##  4 Surface History, Impactors and Climate Evolution

Understanding the surface history of Venus is crucial to constraining any
theory of its long-term climate evolution. Smrekar2018 reviews the literature
on Venus’ internal structure and dynamics. In this section we mostly focus on
implications for the surface features we see today and how those might be
consistent with a hypothesis for the long-term habitability of Venus and a
transition to a more recent ($\sim$1Gyr) hothouse state.

Up to 80% of the Venus surface has volcanic plains and tectonic structures
emplaced over a relatively short geological interval as determined from crater
counts [Ivanov Head (2013), Ivanov Head (2015)]. The cratering record seen in
the plains regions imply surface ages ranging, for example, from $\sim$180Ma
[Bottke . (2016)], to $\sim$300Ma [Strom . (1994)] to $\sim$750Ma [McKinnon .
(1997)]. The relative youth of most of Venus’ surface may be the result of a
large scale lithospheric overturn known as the Global Resurfacing Event (GRE),
or it may be due to the latest GRE in a long sequence of episodic resurfacing
events [<]e.g.,¿Turcotte1993,Strom1994. For example, Kaula1999 constructed a
simple model with outgassing events staggered at time periods of 4.1, 3.8,
3.5, 3.1, 2.6, 2.1, 1.5 and 0.7 Ga constrained by 40Ar measurements. The other
hypothesis for the young surface of Venus is from continuous volcanic
resurfacing [<]e.g.¿[]Basilevsky1997,Bjonnes2012,King2018.

The highland tesserae may be one of the keys to understanding this history.
They are of particular interest because they may contain information about
past crustal differentiation and other processes prior to the loss of any
surface water. Some crater age estimates from the Magellan Mission imply that
the tesserae are $\sim$40% older than the plains [Ivanov Basilevsky (1993),
MS. Gilmore . (1997)]. However, Strom1994 did not agree with this conclusion.
Additional work by Hansen2010 points to the possibility that the Ribbon
Tessera Terrain are older than the surface units identified with the GRE.
Later analysis by IvanovHead2013 implied that tessera are the oldest
stratigraphic unit and that they were created near the beginning of Venus’
surface observable history during the “tectonically dominated regime.”

However, _h_ ow the large basaltic plains were emplaced remains controversial.
A number of authors [<]e.g.,¿[]Herrick1994,Strom1994,BasilevskyHead1996
postulated a nearly global ($\sim$ 80%) geologically instantaneous
(10-100$Myr$) thick ($>$1km) deposition of basaltic material from volcanic
type outflows (GIBVO) that would have buried older craters we cannot observe
today (akin to the GRE mentioned above). The outflow depth requirements are
determined by the size of the largest impact craters that would have to be
completely covered. However, as IvanovHead2013 point out it is possible that
the cratering record previous to GIBVO could have also been erased in some
manner. The GIBVO model was later augmented and became known as the global
stratigraphy hypothesis
[<]e.g.,¿[]BasilevskyHead1996,Basilevsky1997,BasilevskyHead1998,HeadBasilevsky1998.
Yet another hypothesis to explain the Venus surface record was initially put
forward by Phillips1992 and is termed the Equilibrium Resurfacing Model (ERM).
In this model the number of craters observed on Venus today is the result of
an equilibrium between constant crater formation (via impacts) and the removal
of such craters via on-going tectonic or volcanic methods. Monte Carlo
calculations by Bullock1993 and Strom1994 demonstrated why the ERM was not
feasible. Strom1994 decided that the GIBVO was a better fit to their data,
while Bullock1993 preferred a longer timeline of 550$Myr$. More recent Monte
Carlo calculations by Bjonnes2012 show that the ERM is able to fit the
observations.

HansenYoung2007 strove to demonstrate why none of these hypotheses fit all
available observational constraints. HansenYoung2007 then proposed what they
termed the Spatially Isolated Time-Transgressive Equilibrium Resurfacing
(SPITTER) hypothesis to explain more of the observational constraints. It is
not clear that the Venus geological community has settled on any of these
hypotheses. Perhaps one of the largest problems with the global lava
hypothesis is the timescale, volume and depth of the basaltic flows required,
none of which have been observed on any present or previously active volcanic
body in the solar system (including in Earth’s past). The largest known
outflow to date in Earth’s history is the mid-Cretaceous Superplume [Larson
(1991)], which is small by comparison to those envisioned to describe Venus’
resurfacing. At the same time the superplume hypothesis for Venus is
compelling as large amounts of CO2 could have been released at the same time
as the plume event [Caldeira Rampino (1991)]. Large overturn events have been
proposed as an explanation for Venus’ present surface state, but in such a
scenario it is possible to sequester large amounts of CO2 in fresh flood
basalt outflows due to enhanced planetary weatherability
[<]e.g.,¿[]Godderis2003,Cox2016. Large Igneous Provinces (LIPs), on the other
hand, can release copious amounts of CO2 sequestered in some sedimentary
materials [<]e.g.,¿[]Ganino2009 while avoiding the sequestration issues of a
large overturn event. LIPs have been proposed as an explanation for Venus’
present day state as we will discuss below.

Previous simulations by Way2016 showed that Venus could have had temperate
conditions for nearly 2 billion years providing it had a shallow ocean of 310m
in depth, slow rotation rate, and modern orbital elements. Venus might even
have experienced more stable conditions than Earth in its early history since
studies by Correia_Laskar2001 and J_Barnes2016 have shown that low obliquity
states (like that of modern Venus) may be stable over billions of years and we
know that the much shorter Milankovich cycles have had a strong influence on
Earth’s climate through time. Deitrick2018 reviews the influence of such
cycles on the climate of Earth and possible influences on exoplanets. In
addition, Weller2018 has also shown from geological models that early Venus
could have avoided glaciations more easily than early Earth, which experienced
several partial or total snowball periods in its history. If long-term stable
surface conditions are a requirement for life, Venus might have been more
stable and allowed primitive life to fill more ecological niches more quickly
than on Earth. This gives rise to the possibility that life may still exist in
Venus’ upper atmosphere [Limaye . (2018)].

Ernst2017 speculate that “On Venus, voluminous LIP volcanism produced high
levels of CO2 that led to run-away greenhouse effect, and high levels of SO2
that caused acid rain,” but with little supporting evidence.
BullockGrinspoon2001 present a similar hypothesis that involves outgassing of
SO2 and H2O that eventually drive the planet, over 100s of $Myr$, into a
runaway greenhouse state, but do not mention CO2. If Venus had LIP volcanism
then CO2 as well as SO2 can be outgassed if trapped in sediments in the crust
as is seen on Earth
[<]e.g.,¿HeadCoffin1997,Hansen2007,Ernst2017,ErnstYoubi2017. Hence if Venus
had an earlier epoch of liquid water habitability then it is logical to assume
that CO2 would have been trapped in the crust of the planet in the same way it
is trapped on Earth today and LIP volcanism would have been the means to
release that CO2 into the atmosphere.

However, as noted in MacDonald_Wordsworth2017 when the surface temperature is
warmer (T$>$300K, see their Figure 2) more water vapor is injected into the
stratosphere, which stabilizes the lapse rate. Such warm climates (as seen in
the Venus models herein) would prevent the largest plumes from injecting SO2
into the stratosphere, allowing CO2 warming without offsetting cooling by
H2SO4 aerosols.

Another well known mechanism to get Venus from a cool clement state to its
present day hot and dry state was proposed by a number of authors
[<]e.g.¿[]Ingersoll1969,KastingPollack1983,Kasting1984,Kasting1988,TaylorGrinspoon2009
who speculated that water loss via upper atmospheric dissociation and then
hydrogen escape would have eventually made the planet dry. Then, as stated in
TaylorGrinspoon2009 “With the loss of water, the removal mechanism for CO2
would be eliminated, and carbonate rocks on the surface would presumably
eventually be subducted and lost to thermal decomposition, with the CO2 being
irreversibly returned to the atmosphere through outgassing.” This model fits
in with more recent research by Wordsworth2016 who states that the oxygen left
over would eventually find its way to oxidize the mantle and change its redox
state, allowing for enhanced nitrogen outgassing which is compatible with the
nearly 3 bars of N2 we see in Venus’ atmosphere today (also see review by
Lammer2018). However, an alternative hypothesis is proposed by Gillmann2009,
who suggest that the oxidation of the mantle occurred in the first 100 $Myr$
of Venus’ history. They assume the surface was never cool enough to allow
liquid water to condense. The water would again be photodissociated and the
hydrogen would have been lost to space [Lichtenegger . (2016)]. The leftover
oxygen would have dissolved in the magma ocean.

Genda2005 have proposed that the lack of water on Venus and in Venus’
protoplanetary impactors in its early history (in contrast to that of Earth
and its water rich impactors) would explain differences in most of the noble
gas abundances between Venus and Earth because oceanic protoplanets would
enhance atmospheric loss, implying that Venus’ original noble gas abundant
proto-atmosphere survived to present day on Venus, unlike that of Earth. A
lack of water being detrimental to subductive plate tectonics (see Section 8).
Sakuraba2019 have also attempted to get the presently observed nitrogen and
noble gas abundances via impact degassing and atmospheric erosion [<]also see
work by,¿[]Pham2011, but unlike Genda2005 they believe late accretion may have
further influenced the atmosphere of Venus.

More recent work by Gillmann2016 show that large impactors (400-800km in
diameter) can cause atmospheric erosion and escape and deposit energy in the
crust and mantle. They believe the latter can cause a thermal anomaly in the
crust and mantle triggering large scale volcanic events at the impact region
and the antipode. This in turn may deplete the upper mantle of volatiles and
lead to water loss in the early atmosphere, or conversely provide a volatile
heavy atmosphere with extreme temperatures for billions of years. In a sense
this is similar to a theory by Davies2008 who propose a mega-collision (akin
to that of the Earth’s moon-forming impact) to dry out the interior of the
planet. But thus far no large Venus impactor simulations have been utilized to
examine such a scenario, as has been done for Earth’s moon-forming collision
[<]e.g.,¿[]Canup2004.

To summarize, a number of mechanisms exist by which early Venus could have
condensed liquid water on its surface. The key ingredient is that it must have
been cool enough for long enough in its early history. As shown by
Yang2014,Way2016,Way2018, the rotation rate of a planet greatly affects its
climate dynamics. Specifically, for very slow rotation a large contiguous
water cloud forms at the substellar point, increasing the Bond albedo markedly
and keeping surface temperatures moderate for insolation values up to nearly
three times that of modern Earth’s 1361 W m-2. In Paper I Way2016 we
demonstrated that early Venus could have had consistently habitable conditions
throughout its early history if it began with sufficiently slow rotation. In
the next section we review what is understood about the possible evolution of
Venus’ spin-orbit state.

In our scenario, early Venus’ has the earliest consistent liquid water
habitability in the solar system followed by Earth and then Mars. This is a
broader statement of the Faint Young Sun Paradox (FYSP), the challenge of
explaining how early Earth, not to mention Mars, could have been warm and wet
early in their histories when the Sun was 25-30% dimmer than today
[<]e.g.,¿Feulner2012. There is still debate in the ancient Earth GCM community
about the actual composition and thus temperature of early Earth’s atmosphere
given observational proxies for CO2 that span orders of magnitude, though
models suggest that the range encompasses several viable scenarios
[<]e.g.,¿[]Charnay2013,Wolf_Toon2013,Kunze2014,LeHir2014,Charnay2017,Krissansen-
Totton2018. These GCM studies and most proxies [<]e.g.,¿[]Spencer2019 are from
the Archean rather than the late Hadean, but there is some evidence that
habitable surface conditions existed well back into the Hadean
[<]e.g.,¿[]Harrison2009,Arndt2012.

It is interesting to note that recent atmospheric pressure proxies from the
late Archean imply an atmospheric pressure less than half that of today [D.
Som S. Catling . (2012), S. Som . (2016)]. Atmospheres thinner than modern
Earth’s are less likely to avoid snowball conditions, yet the literature above
notes that there is geological evidence that Earth was not in a snowball state
during much of the late Archean that the pressure proxies correspond to.
Regardless, for this reason we feel is it necessary to explore the
possibilities of lower atmospheric surface pressures in Venus’ climatic
history as described for Simulations 26-30 in Section 6.

The FYSP for Mars remains difficult to resolve [<]e.g.,¿[]Wordsworth2016AREPS
partly due to the fact that 3-D GCMs have traditionally struggled to
consistently sustain large-area liquid water conditions over millions of years
[<]e.g.,¿[]Goldblatt2009,Kasting2010,Kienert2012,Feulner2012,Haqq-Misra2008
without snowball type conditions. Long-standing solutions involving large
amounts of atmospheric CO2 are inconsistent with unobserved carbonate deposits
expected from such CO2 dominated atmospheres [Shaw (2018)] and are
insufficient in isolation to produce above-freezing conditions. One possible
solution to the lack iof surface carbonates was proposed by Kasting2012. Other
solutions to Mars’ FYSP exist that involve H2 with CO2 as the background gas
[<]e.g.¿Wordsworth2017,Ramirez2014b,Ramirez2018,Haberle2019, although
presently there appears to be little consensus in the community.

##  5 Rotation and Obliquity evolution

To the best of our abilities we would like to constrain the obliquity and
rotational history of Venus to better constrain these important inputs for
climate models. This is limited by the absence of any direct information about
Venus’ initial rotation and obliquity and the fact that impacts likely play a
significant role in the early rotational history of the terrestrial planets
[<]e.g.,¿[]LissauerKary1991,DonesTremaine1993. On Earth a variety of means
exist to obtain some constraints using dynamical modeling combined with
geological data when available
[<]e.g.,¿[]Hays1976,Park1987,Imbrie1992,Matthews1997,Petit1999,Palike2000,Palike2004,Olsen2019
and there has been modest success doing the same for Mars
[<]e.g.,¿[]Cutts1982,Laskar2002Mars,Laskar2004,Byrne2009,Dickson2015,Bierson2016.
For Earth, an additional constraint is provided by the Moon, which has
predictably affected the evolution of Earth’s rotation and damped obliquity
excursions over its history [Zahnle Walker (1987), Lissauer Chambers (2011)].
However, until and unless geological observables become available to constrain
dynamical models, only plausible scenarios for the rotational and obliquity
history of Venus can be defined.

vanHoolst2015 summarizes much of the literature on the rotational evolution of
Venus throughout its history. We summarize some of the work on this subject
below and add some additional estimates. First we look at the history of
studies of the possible spin evolution of Venus.

In the 1960-70s several authors investigated the possibility that Venus’
rotation period was correlated with its synodic period [Goldreich Peale
(1966), Gold Soter (1969), Gold Soter (1979)]. Goldreich1966 states, “the
presence of the Earth may have stabilized the sidereal rotation period of
Venus at the value of 243.16 days retrograde.” An equilibrium between the
atmospheric and body tide of Venus was first proposed by Gold_Soter1969 to
explain Venus’ non-synchronous rotation period, based on the incorrect belief
at that time that Venus always showed the same face at each inferior
conjunction with Earth as proposed by Goldreich1966.

The first work to analytically look at Venus’ rotation rate and the role of
atmospheric tides was by Ingersoll1978 who extended the earlier work of Lord
Kelvin Thompson1882, ChapmanLindzen1970 and Munk_MacDonald1960. They mention
that “Venus probably originated with a retrograde rotation in order to have
evolved to the current retrograde state.” In the 1980s this work was further
extended in a series of papers [Dobrovolskis Ingersoll (1980), Dobrovolskis
(1980), Dobrovolskis (1983)]. It was clear that Venus’ rotation rate was
probably determined by a balance between the solid body tidal dissipation and
the thermal tides of its thick atmosphere with the sun. Core-mantle friction
(CMF) can also play an important role in slowing the spin rate of Venus, as
first explored by Goldreich_Peale1970. Goldreich_Peale1970 were also the first
to demonstrate that core-mantle viscous coupling can drive the obliquity to 0∘
when less than 90∘ and to 180∘ if it is greater than 90∘ over time.

This remained the state of understanding of Venus’ rotational history until
the early 2000s when the long-term evolution of its spin state of Venus was
investigated in a series of papers by
Correia_Laskar2001,Correia2003a,Correia2003b, who suggested that Venus may
have rotated faster in the past, and possibly prograde. It also became clear
that at faster spin rates CMF plays an important role in slowing the rotation
of the planet, but less so at slower spin rates.

Once a planet is spinning more slowly CMF may play an important role in
obliquity variations [<]e.g.,¿[]Correia2003a. Correia_Laskar2001 explored a
number of stable obliquity and spin states of Venus while more recent work by
J_Barnes2016 has investigated how stable the obliquity of Venus might be
though time.

The work on the thermal tides of Venus had led researchers to assume that its
effects would be minor (as it is for Earth) for atmospheres of modest density
(e.g. 1 bar). However, more recent work by Leconte2015 has demonstrated that
thermal tides arising from even 1 bar atmospheres can be significant depending
on the distance to the host star and the host star’s mass. Leconte2015 show
that even if modern Venus had a 1 bar atmosphere the tidal torques would still
be quite significant.

Barnes2017 used an equilibrium tide model with a constant phase lag (CPL) to
find that Earth could have ended up tidally locked today (after 4.5Gyr) had it
started with a rotation rate of 3 Earth Days or longer (the latter more likely
if Earth had no satellite). We have applied the same Equilibrium Tide Model
(EqTide https://github.com/RoryBarnes/EqTide) from Barnes2017 to Venus to
explore how long it would take Venus to reach a tidally locked state only from
solid body tides. As shown in Figure 1 using CPL theory we find that Venus
could have been tidally locked within 684$Myr$ if it started with a prograde
rotation period of 3 Earth days and zero obliquity. Unfortunately the EqTide
model we utilize does not support retrograde spin states, but we expect the
differences to be minor. We will continue to explore these issues in a future
work using the simulator vplanet [R. Barnes . (2019)] once this functionality
is added. Figure 1 gives further examples for CPL and Constant Time Lag (CTL)
theory results using EqTide. For input parameters we assume that the tidal
dissipation factor Q=12 and Love number of degree 2 k2=0.3. These are the same
numbers used for the modern Earth in Barnes2017. Recent work by Henning2014
demonstrates that our choice for Q may not be unreasonable for Venus.
Henning2014 give estimates of Q for Earth-like planets (see their Fig 15, top-
center-row plot) with orbital periods from 0 to 200 days. Venus’ 224 d period
is slightly outside the range they explore (but can be anticipated from the
trend visible in their figure). Our assumption of Q=12 is not far off the
Henning2014 ‘Warm Earth 2’ estimate in their Fig 15. As an aside, Q and k2 are
poorly constrained for present-day Venus. We have even fewer constraints on
these values for an ancient Venus, but perhaps those values would be more
Earth-like than present day Venus. For example, present day Venus’ time lag
may not be the same as Earth’s because of higher internal temperatures
[MacDonald (1962), Henning Hurford (2014)]. Historically Goldreich_Soter1966
estimated that Q$<$17 for Venus, Lago1979 had values up to Q$\sim$40 while
Leconte2018 estimate Q$\sim$100.

More recently Venus’ tidal love number was estimated by Konopliv1996 using
Magellan and Pioneer Venus Orbiter data to be k2=0.295$\pm$0.066 implying the
core is liquid [Yoder (1997)]. Work by Zhang1992,Xia2002 have estimated
k2=0.18 $\sim$ 0.26. A smaller value (k2=0.17) would imply a solidified iron
core which is not consistent with Konopliv1996. Modeling work by Dumoulin2017
are consistent with the work of Konopliv1996 as well as our own modeling
choices (discussed above) of Q=12 and Love number of degree 2 k2=0.299 (see
Table 3 in Dumoulin2017). Regardless, if one uses higher values of Q and/or
lower values of k2 for ancient Venus it is sufficient to say that equilibrium
tide theory predicts that the CPL and CTL for Venus estimates for tidal
locking will be longer than those presented in Figure 1. The values in Figure
1 then represent _lower limits_ to tidal locking for a given starting
rotational period. As a caveat there is a debate in the dynamics community
about the appropriateness of the CPL and CTL approaches [Efroimsky Williams
(2009), Efroimsky Makarov (2013), Touma Wisdom (1994), Greenberg (2009)]. so
these tidal locking timescales should be viewed with some caution in the
context of the CPL and CTL models used herein. Of course we do not take into
account magnetic braking to see how the Sun’s natural spin-down might affect
the tidal evolution of Venus, nor do we assume that Venus’ orbital
characteristics would have changed over the timescale of our calculations, the
latter being one of the criticism when applying CPL/CTL to evolving systems
[Efroimsky Williams (2009), Efroimsky Makarov (2013)].

Figure 1: Venus tidal locking timescales using equilibrium tide constant phase
(a) and time (b) lag theory. Note that the y-axes have different limits.

Recent work [Green . (2019)] has investigated the influence of a hypothetical
shallow ocean on Venus (water equivalent layers of $\sim$330 meters deep and
830m) using present day topography and a range of initial rotation periods.
The most dissipative scenario predicts a slow down of 72 days per million
years. The latter result may be surprising until one recalls that tidal
dissipation in Earth’s oceans is larger than that of the Earth’s solid body
tides [<]e.g.,¿[]Munk_MacDonald1960.

It is likely that Venus was initially a prograde spinning body like the other
3 terrestrial planets in our solar system. The prograde hypothesis goes back
at least to [Alfvén (1964)] and more recent work by
LissauerKary1991,DonesTremaine1993,KaryLissauer1995 would also support the
idea of a primordial prograde Venus, barring the effects of a late large
impactor as discussed in those works and that below. From that starting point
we find the following perhaps the most compelling answer to Venus’ present day
spin state. As shown above there are models that can drive the planet toward a
tidally locked state rather quickly. Core-mantle friction damps obliquity
perturbations which drive the spin rate to sub-synchronous prograde rotations.
Then atmospheric tides would reverse the spin to a retrograde equilibrium.
These atmospheric tides continue to prevent the planet from being tidally
locked and that is the state the planet has been in since that time as shown
in some of the work of Correia_Laskar2001.

Yet it has long been speculated that Venus’ current retrograde rotation state
is the result of a large impactor early it its history [McCord (1968), Singer
(1970), French Singer (1971), Counselman (1973), Burns (1973), Ward Reid
(1973), Harris (1978), Alemi Stevenson (2006), Davies (2008)]. The large
impactor hypothesis may also explain a possibly very dry Venus interior, a
lack of oxygen in the atmosphere and reconcile the 40Ar results that imply it
is less degassed than Earth [Davies (2008)]. However, as pointed out in
WardReid1973 it is possible for a small impactor (less than 1% of the Moon’s
mass) to drive Venus retrograde if the planet has already spun down
considerably due to tidal dissipation with the Sun (Figure 1).

Unfortunately, there is little hope that we will ever truly know the rotation
rate of Venus through time without a way to either measure its “Geological
Orrery” as on Earth [Olsen . (2019)] and tie that to dynamical models, or find
evidence that an impactor played a role in its rotational and hence
geochemical evolution as seen in investigations on Earth related to late
accretion and its effect on different isotope abundances [<]e.g.,¿[]Varas-
Reus2019. In fact Brasser2016 and Mojzsis2019 prefer the hypothesis that the
Earth’s Late Veneer was mainly delivered by a single Pluto- or Ceres-sized
impactor. Hence if a larger object was involved in the late evolution of
Venus’ spin or obliquity it may be possible to detect its geochemical
fingerprints in a future in-situ mission.

For _most_ of our climate simulations we assume, rightly or wrongly as
discussed above, that Venus has had the same retrograde rotation and its
present day obliquity for the past 4.2$\times$109 years. We have included a
few faster rotation rates that approach the ‘fast rotators’ described in
Yang2014,Way2018, meaning sidereal day lengths of 16 and 64 times that of
present day Earth.

##  6 Methods

All our simulations use ROCKE-3D [Way . (2017)] a three-dimensional (3-D)
General Circulation Model (GCM) developed at the NASA Goddard Institute for
Space Studies (GISS). Radiative transfer in ROCKE-3D uses k-coefficients
derived from the HITRAN 2012 line list, as well as the MT-CKD 3.0 water vapor
continuum and CO2 collisionally-induced absorption and sub-Lorentzian line
shapes, as described in references cited in [Del Genio, Way . (2019)]. HITRAN
2012 is accurate for temperatures below 350 K, as shown by [Kopparapu .
(2013)]. As we will discuss below, a small number of our simulations exhibit a
growing radiation imbalance with time, indicative of continually rising
temperatures and a transition to a runaway greenhouse. If our objective were
to determine the threshold for a runaway, we would need to use the more
comprehensive HITEMP line list for these simulations, as discussed by
[Kopparapu . (2013)]. Our purpose, though, is simply to identify such cases
and exclude them from further analysis.

Most simulations use modern Venus’ current orbital parameters, slow retrograde
rotation period (-243 Earth sidereal days in length) and orbital period (224
sidereal days). In Paper I (plotted herein with ID = B) we looked at a faster
rotation period (16 x modern Earth’s sidereal day length) to see how the
planet’s early climate might have responded, and we also look at 16 & 64 day
retrograde rotation periods in Venus’ early history in this work. Our focus is
on changing insolation, topography, land/sea mask, surface water availability
and atmospheric constituents. We motivate our choices below.

All simulations discussed below are outlined in Table 1. Our focus is on the
retention and stability of surface liquid water on Venus over time. Therefore
we simulate four types of planets with surfaces that differ in the amount of
water they contain and how that water is allowed to interact with the
atmosphere. ROCKE-3D allows for 3 types of surface water: Soil moisture at and
beneath the surface with no standing bodies of water; “dynamic” lakes whose
depth and area vary with time and that can appear or disappear as the
competition between precipitation and evaporation dictates; and deeper oceans
with permanent boundaries and an effectively infinite source of water for the
atmosphere. All planets with oceans are fully dynamic. For more details on the
capabilities of such oceans see Way2017.

  * •

Arid Venus: This planet has modern Venus topography, but only contains 20cm of
water in the subsurface soil layers, soil consisting of 100% sand, and no
surface standing water at the start of the simulation. The atmosphere is
initialized with zero water vapor and an isothermal temperature profile at
300K. This initial condition is similar to that of Kodama2019,Abe2011 who
attempt to limit the amount of water vapor in the atmosphere (a strong
greenhouse gas) and subsequently push the inner edge of the habitable zone
farther inward. However, Kodama2019,Abe2011 use modern Earth’s rotation rate
for all their experiments.

  * •

10m-Venus: Uses modern Venus topography and places a 10 meter liquid water-
equivalent layer in the lowest lying topographic areas. These are treated by
the model as lakes, which have no circulation. The soil is a 50/50 sand/clay
mix as used in Yang2014,Way2016,Way2018.

  * •

310m-Venus: Similar to 10m-Venus, except with a 310 meter water equivalent
layer again spread in the lowest lying regions. This is the same topography
used in Way2016 simulations A,B and D.

  * •

158m-Aqua: This is a simple aquaplanet configuration that is commonly used in
the exoplanet community. It uses a fixed 158 meter deep ocean, which
corresponds to the bottom of the fifth layer of the ROCKE-3D ocean model. It
is a bit shallower than the mean depth of the 310m-Venus ocean, and therefore
comes into equilibrium a bit faster while still having a similar heat
capacity, while including horizontal heat transport as well as wind-driven and
thermohaline overturning circulations.

  * •

310m-Earth: Similar to 310m-Venus, but using a modern Earth-like land/sea mask
with a 310m deep bathtub dynamic ocean (i.e., every ocean grid cell is of a
fixed depth of 310m). We call this an Earth-like land/sea mask since it is not
exactly modern Earth, but has some modest changes as shown in Way2018 Figure
8.

The five planets above are then given four types of atmospheres and four
different insolations as described below:

  * •

Simulations 1-5: These have a 10 bar 100% CO2 atmosphere using a solar
spectrum and insolation from 4.2Ga from the work of Claire2012. CO2 was
probably the dominant gas in Earth’s early atmospheric evolution
[<]e.g.,¿[]Kasting1993a. We pick atmospheric pressures of 1 bar (see next
bullet point) and 10 bar to cover the Kasting1993a ranges (see their Fig 2).
Our 10 bar results either equilibrate at a temperature beyond that at which
our radiative transfer is accurate, or do not reach equilibrium and the
temperatures attained at the time the experiments were terminated are already
beyond the upper limits of our radiation tables. We report the results of
these experiments in Table 2 below simply as a guide for future research, but
we exclude them from our analysis in Figs. 2-8. The 10 bar simulations use a
modern Venus rotation rate and obliquity.

  * •

Simulations 6-10: Similar to Simulations 1-5, but these use a 1 bar 97% CO2
and 3% N2 atmosphere at 4.2Ga.

  * •

Simulations 11-15: As in Simulations 6-10 but with a rotation period of –16
sidereal Earth days to place the planet on the edge of the fast rotator regime
as described in Yang2014,Way2018. This allows us to explore the possibility
that the planet was rotating more quickly in its early history than today. The
choice of a retrograde rotation rate was chosen to be consistent with the
present day retrograde rotation, but unpublished simulations with prograde
rotation rates with these values produce very similar temperatures. Note that
the work of Correia_Laskar2001,Correia2003b indicate that prograde rotation
rates of 16 days for Venus put its spin axis (obliquity) in a possibly chaotic
regime. However, other work by J_Barnes2016 indicate that low obliquity
retrograde rotation rates generally have more stable spin axes. Even if there
are spin axis variations on geological timescales, it is not possible for us
to model those here given that ROCKE-3D simulations are limited to $<$ 10,000
years in length.

  * •

Simulations 16-20: As in Simulations 6-10 but with a rotation period of –64
sidereal Earth days. This allow us to explore the possibility that the planet
was rotating somewhat more quickly in its early history but still in the
slowly rotating dynamical regime. Again, prograde rotation rates were also
used in unpublished results and have similar global surface temperature
values.

  * •

Simulations 21-25: These simulations use an atmospheric composition and
pressure very similar to modern Earth, namely an N2-dominated atmosphere with
400ppmv CO2 and 1ppmv CH4 with a 1013mb surface pressure. They also use a
solar spectrum and insolation at 2.9Ga from Claire2012. The rotation rate is
the same as modern Venus.

  * •

Simulations 26-30: Similar to Simulations 21-25, but with a lower atmospheric
surface pressure of 250mb. This is again in the interest of comparative
climatology since the Archean atmospheric pressure proxy work of
Som2012,Som2016 suggests that Earth may have had a surface pressure similar to
250mb at this time.

  * •

Simulations 31-35: Similar to Simulations 21-25, but now using a solar
spectrum and insolation from 0.715Ga from the work of Claire2012.

  * •

Simulations 36-40: Again, similar to Simulations 21-25, but now using a modern
solar spectrum and insolation.

  * •

Simulations 41-45: Similar to Simulations 21-25, but now using a modern solar
spectrum, but with insolation set to 1.26 times Venus’ present day insolation
(2.4 times modern day Earth’s insolation) to test the boundaries of the inner
edge of the habitable zone as in Way2018. Two of these simulations are also
out of radiation balance and trending toward a runaway greenhouse state, and
thus we do not analyze them further.

Table 1: Experiments

ID |  Topographya | Epoch |  Insolationb |  Pc | Spin |  N2 |  CO2 |  CH4 | Soil
---|---|---|---|---|---|---|---|---|---
|  | Ga |  S0X/W m-2 | bar | days | ppmv | ppmv | ppmv |  Typed
01 | Arid-Venus | 4.2 | 1.396/1913.6 | 10 | -243 | 0 | 1000000 | 0 | S
02 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
03 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
04 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
05 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
06 | Arid-Venus | 4.2 | 1.396/1913.6 | 1 | -16 | 43000 | 970000 | 0 | S
07 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
08 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
09 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
10 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
11 | Arid-Venus | 4.2 | 1.396/1913.6 | 1 | -64 | 43000 | 970000 | 0 | S
12 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
13 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
14 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
15 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
16 | Arid-Venus | 4.2 | 1.396/1913.6 | 1 | -243 | 43000 | 970000 | 0 | S
17 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
18 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
19 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
20 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
21 | Arid-Venus | 2.9 | 1.47/2001.0 | 1 | -243 | 1012599 | 400 | 1 | S
22 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
23 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
24 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
25 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
26 | Arid-Venus | 2.9 | 1.47/2001.0 | 0.25 | -243 | 1012599 | 400 | 1 | S
27 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
28 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
29 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
30 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
31 | Arid-Venus | 0.715 | 1.71/2358.9 | 1 | -243 | 1012599 | 400 | 1 | S
32 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
33 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
34 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
35 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
36 | Arid-Venus | 0.0 | 1.9/2601.0 | 1 | -243 | 1012599 | 400 | 1 | S
37 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
38 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
39 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
40 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
41 | Arid-Venus | Future | 2.4/3266.0 | 1 | -243 | 1012599 | 400 | 1 | S
42 | 10m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
43 | 310m-Venus | ” | ” | ” | ” | ” | ” | ” | S/C
44 | 158m-Aqua | ” | ” | ” | ” | ” | ” | ” | -
45 | 310m-Earth | ” | ” | ” | ” | ” | ” | ” | S/C
Dd | 310m-Venus | 2.9 | 1.47/2001.0 | 1 | -16 | 1012599 | 400 | 1 | S/C

aTopography: Arid-Venus=Only Ground Water, no surficial reservoirs, 20cm water
in soil, with modern Venus topography; 10m-Venus=10m Water Equivalent Layer
(WEL) spread in lowest elevations as lakes with modern Venus topography;
310m-Venus=310m deep ocean with modern Venus topography; 310m-Earth=Modern
Earth-like topography with 310m deep ocean; 158m-Aqua=158m deep aquaplanet.
bInsolation: S0X = multiple of amount that Earth receives today in insolation
(S0=1361 W m-1).
c Pressure in bar. d S=100% Sand, S/C=50/50% Sand/Clay, - = Not Applicable,
100% ocean.
d Simulation D from Paper 1 [Way . (2016)]. Most similar to ID 8.

Most simulations except 158m-Aqua use a fixed ground albedo of 0.2 (thermal
conductivity = 0.26 W m-1 K-1) and 50/50 mix of sand/clay soil following the
work of Yang2014,Way2016,Way2018. The Arid-Venus simulation uses the same
albedo (0.2), but utilizes a 100% sand soil, rather than the sand/clay mix in
other simulations. The advantage of using sand is that it more quickly loses
and absorbs water. This allows the ground hydrology to come into balance more
quickly than other soil types. This is because in the Arid-Venus simulations
we are focused on water availability to/from the atmosphere from/to the soil
and hence the amount of total water vapor acting as a greenhouse gas in the
atmosphere.

##  7 Results and Discussion

Simulations 1-5, all with a 10 bar pure CO2 atmosphere and 4.2 Ga insolation,
are uniformly uninhabitable regardless of the surface water reservoir and
topography (see Table 2 in the Supplementary Information). The driest planet
(Arid Venus) does reach equilibrium, but with a surface temperature of 262∘C,
well above the accuracy limits of the radiation parameterization used by
ROCKE-3D. The other four planets are also well above 100∘C at the point at
which they were terminated and are not converging to equilibrium. Given the
greater water reservoirs in these simulations, they are likely to be
approaching a runaway greenhouse state.

In Figure 2 we show several different possible evolutionary scenarios for
Venus derived from the other experiments in Table 1. In all such scenarios we
assume that Venus had surface liquid water in varying amounts at model start,
as described in Section 6. The colors in this figure differentiate groups of
simulations with different insolation, rotation, surface pressure, and/or
atmospheric compositions, while the numbers 1-5 and corresponding symbols for
each color delineate the range of climates obtained for different surface
water reservoir and topography assumptions.

Figure 2: Atmospheric evolutionary scenarios for Venus. Top x-axis is
insolation relative to modern Earth (1.4 = 1.4 $\times$ 1361 W m-2). Note that
the gray, magenta and red data in the left-most section of this plot are all
for 4.2 Ga simulations. They are separated purely for visual effect and do not
reflect differences in epoch or insolation. The same is true for the cyan, and
orange data, all of which correspond to 2.9 Ga. The numbers oriented
vertically along each set of simulations correspond to the different water
reservoir/topography types: 1=Arid-Venus, 2=10m-Venus, 3=310m-Venus
4=158m-Aqua, 5=310m-Earth. See Table 2. Figure 3: Shortwave cloud radiative
forcing for the same evolutionary scenarios, defined as the difference between
the solar radiation actually absorbed by the planet and how much would be
absorbed if clouds were transparent. This is an estimate of how effective the
clouds are at shielding the planet from the star’s intense radiation. Figure
4: Percentage areal coverage of high level clouds (PCLDH). In Way2018 it was
shown that PCLDH plays a key role in the cloud albedo feedback for slow
rotators.

###  7.1 4.2Ga

Since there are major uncertainties about what Venus’ initial rotation rate
was (See Section 5) we explored early post-magma ocean scenarios at 4.2Ga with
three different retrograde initial rotation periods in the left hand part of
Figures 2,3,4: –16 days (gray; experiments 6-10), –64 days (magenta;
experiments 11-15) and –243 days (red; experiments 16-20). Each assumes a 1
bar CO2-dominated atmosphere. As one would expect from the studies of
Yang2014,Way2018 the faster spin rate simulations generally have higher
temperatures because of the cloud processes discussed in those papers, but
almost all of them reach equilibrium at a habitable global mean surface
temperature. However, the clouds also differ to some degree because of water
availability. Contrary to the work of Abe2011,Kodama2019 the Arid-Venus cases
all have higher surface temperatures than their counterparts. This is because
those previous works used modern Earth’s rotation rate, whereas the cloud
processes on these slower rotating worlds better regulate the climate, more so
the more water that is available for cloud formation. This analysis is backed
up by Figure 3 where we plot the shortwave cloud radiative forcing (SWCRF).
The Arid-Venus simulations have the smallest (in magnitude) values, because a
drier planet has less reflective clouds with less condensed water. In Figure 4
we show the percentage of high level clouds (PCLDH), the dominant of the three
cloud types (high, medium, low) in Table 2. Here the distinction between the
Arid-Venus simulations and the others is not consistent across the different
rotation periods, suggesting that in some cases middle and/or low level clouds
make important contributions to SWCRF.

###  7.2 2.9Ga

Here we plot two different sets of simulations for N2-dominated atmospheres: 1
bar (cyan; simulations 21-25) and 250mb (orange; simulations 26-30). In effect
these portray representative possible atmospheres for an ancient Venus with
liquid water that has evolved from an early CO2-dominated atmosphere to a more
Earth-like composition via the carbonate-silicate cycle feedback that is
believed to regulate CO2 on planets with liquid water. In both cases the
Aqua-158m simulations have the highest mean surface temperatures with the
Arid-Venus a close second, but all 10 simulations have moderate surface
temperatures fairly similar to modern Earth. However as for simulations 6-20
the shortwave cloud radiative forcing is again the smallest for the Arid-Venus
simulations (Figure 3) while also having less high cloud in (Figure 4) than
the simulations with more surface water. Unsurprisingly, the thin 250mb
atmospheres (simulations 26-30) have cooler surface temperatures in Figure 2.
Simulations 27 & 28 have lower mean surface temperatures than modern Earth.
The surface temperature field for simulation 28 is plotted in Figure 5 for
reference. It exhibits fairly uniformly warm oceans, a signature of slowly
rotating planets. Continental temperatures are cooler, slightly below freezing
on average, due to nighttime cooling that offsets daytime warming. We also
plot Simulation D from Paper I. The other 3 simulations from Paper I have
similar values to their corresponding simulations herein. Simulation D is an
N2-dominated atmosphere, but is otherwise similar to ID 8 (a CO2 dominated
atmosphere) in Tables 1 and 2. It has a lower mean surface temperature than ID
8 as expected, but is significantly higher than the other simulations with
larger rotation periods (ID=21–30) at 2.9Ga. It has less short wave cloud
radiative forcing (Figure 3) and lower percentage of high level clouds (Figure
4) compared to the other 310m-Venus simulations at 2.9Ga. This is expected
given its faster rotation period, stronger Coriolis force, and less contiguous
clouds at the substellar point as discussed in Paper I.

Figure 5: Simulation 28: mean surface temperature over 1/6th of a diurnal
cycle. The substellar point is centered over the middle of the plot.

###  7.3 0.715Ga

This epoch captures a possible final habitable phase on Venus, if the thick
CO2 atmosphere we see today was created by volcanic emissions during the
global resurfacing event(s). The spread in surface temperatures between
simulations remains about the same as in previous epochs, but for all surface
types the temperature is warmer than at 2.9 Ga because of the brighter Sun.
Again, the Aqua-158m is has the highest surface temperatures with the Arid-
Venus close behind as in the 2.9Ga epoch. The SWCRF is somewhat larger in
general, due to both the stronger insolation and slightly reduced high cloud,
but again the Arid-Venus has the smaller effect. The spread in high level
clouds has shrunk considerably from the 1 bar simulations at 2.9Ga.

###  7.4 Present Day

This suite of simulations at Venus’ present day insolation are designed to
demonstrate that even under today’s Sun the slow rotation cloud feedback
effect would have remained strong as Venus’ atmospheric pressure and
composition remain unchanged. This points to the idea that it was not an
increase in insolation that drastically changed Venus’ clement climate of
earlier epochs, but rather something else, which we speculate to be
multiple/simultaneous large igneous provinces. The Arid-Venus simulations
again have the highest temperatures and corresponding smallest SWCRF. This is
more along lines of what we saw with the simulations at 4.2Ga, with climate
forcing by a stronger Sun replacing climate forcing by a thicker greenhouse
gas atmosphere as the primary reason for a warm climate.

###  7.5 Future

Our last set of simulations at insolation values 2.4 times that of present day
Earth are meant to show how long a temperate Venus-like world could have
remained habitable for a given surface type. Our 10m-Venus and 310m-Earth
simulations are not in equilibrium and so are not plotted. The Arid-Venus
simulation is at nearly 100∘ Celsius. It appears to be approaching radiative
equilibrium, but the simulation crashed after 20 years so it is difficult to
be certain. At this point the cloud/albedo feedback for the Arid-Venus case
has decreased to Earth-like values (the SWCRF is a mere –50 W m-2). This
simulation has the lowest value of PCLDH, which makes it hard to counter the
increased insolation at this time in order to keep mean surface temperatures
below the boiling point of water.

###  7.6 General Trends

A few relatively consistent trends are apparent from our simulations. First,
the Arid-Venus simulations tend to have the highest surface temperatures,
smallest values of SWCRF and lowest percentages of PCLDH. In many cases the
10m-Venus simulations are next, but not always. This may point to the fact
that even with 10m of available water the cloud/albedo feedback is generally
effective at shielding this slowly rotating world from the intense solar
radiation at all epochs when considering what modern Earth receives. The
simulations with generally higher water availability and similar percentages
of land-to-sea (310m-Venus and 310m-Earth) tend to cluster together in Figures
2,3,4. The 158m-Aqua simulations seem to float in between, perhaps because of
the lower surface albedo in combination with the cloud albedo feedback.

Figure 6: Upper panel: Bond albedo vs. incident solar flux for the planets in
Figure 2. Middle panel: Predicted versus actual Bond albedos for the same set
of planets after DelGenio2019. Lower panel: Predicted vs. observed surface
temperatures after DelGenio2019.

Even without a transition to a runaway greenhouse, prior water loss due to the
onset of a “moist greenhouse” state may have been important to Venus’
evolution. A number of simulations in Table 2 contain stratospheric water
concentrations (Qtop) greater than 3$\times$10-3 kg kg-1, the traditional
Kasting1993b limit for onset of the moist greenhouse. However, recent work by
Chen2019 for M-star planets has demonstrated that previous work may have
overestimated water loss rates. Hence we should exercise more caution in using
3$\times$10-3 kg kg-1 as a hard value for the moist greenhouse until similar
models are applied to G-star planets. Column Qsurf is a check on the amount of
water vapor at the surface of the model. ROCKE-3D runs with a fixed molecular
mass at model start and ignores the spatially/temporally variable mass of
water in calculating pressure gradients, so it is important to keep track of
whether water becomes a non-negligible fraction of atmospheric mass (e.g., 20%
of the total or more) as the dynamics in the model will begin to be outside an
acceptable range. Only in simulation 45 does this value go over the 20% limit.

Figure 6 shows calculations for what exoplanet astronomers might find for a
population of “exo-Venuses,” some of them habitable and some not, in future
observations. We use an ensemble of ROCKE-3D simulations of a variety of rocky
planet types from which predictors for Bond albedo and surface temperature
have been derived using insolation and star temperature as inputs
[<]¿[hereafter DG19]DelGenio2019. Figure 6 applies the predictor to our Venus
evolutionary scenarios to determine the predictability of albedo and surface
temperature. In general Bond albedo increases with insolation in the Venus
simulations (upper panel), the exceptions being 4 of the Arid-Venus cases with
limited surface water, fewer clouds and thus lower albedos than our other
simulations. The DG19 predictor (middle panel) works well for all but these 4
cases, since it predicts a high cloud-controlled albedo for the wetter planets
that have more and/or thicker clouds than others. For surface temperature
(bottom panel), the predictor tends to underestimate the actual temperatures
by roughly 20∘ or less in most cases, but by up to 50∘ for the hotter, drier,
marginally habitable Arid-Venus cases.

Figure 7: GWTR left panels, LakeFR right panels for Top to Bottom: Simulation
21 and 26 (Arid-Venus). Figure 8: GWTR left panels, LakeFR right panels for
Top to Bottom: Simulation 22 and 27 (10m-Venus).

Figure 7 shows (left panels) the vertically integrated soil moisture and
(right panels) lake fraction for two of our Arid-Venus simulations at 2.9 Ga
with Earth-like atmospheres: Experiment 21 (1 bar, top panels) and Experiment
26 (250 mb, bottom panels). The Arid-Venus cases are of particular interest
because they are initialized with a spatially uniform subsurface soil water
reservoir and no standing water bodies, and they then equilibrate to a
heterogeneous distribution of surface and subsurface water depending on the
climate and circulation and thus the local precipitation-evaporation
competition. The 1 bar atmosphere is typical of the behavior of most of the
experiments in Table 1: In equilibrium, subsurface water collects primarily in
the equatorial region where rising motion and precipitation is prevalent
during the daytime, and especially in highland regions such as Aphrodite
Terra. Lakes (which are not present in the initial condition) also form
preferentially over the equatorial highlands. The 250 mbar simulation is an
outlier, with subsurface water and lakes arising primarily at high latitudes.
This is reminiscent of what is observed for the methane-ethane lakes on Titan,
which Mitchell2008 is able to reproduce in a GCM when a limited subsurface
methane reservoir is assumed.

Figure 8 shows the same quantities for the analogous 10-m Venus simulations
(Experiments 22 and 27). These differ from the Arid-Venus cases not only
because the planet contains more water, much of it in surface lakes rather
than subsurface soil moisture, but also because the lakes fill the lowlands at
the start of each simulation rather than being distributed uniformly across
the planet. With a larger water reservoir than that for the Arid-Venus
planets, soil moisture collects throughout the tropics in the equilibrated
climate, but still with a slight preference for the highland regions even
though there is no standing water in the highlands in the initial condition.
But unlike the Arid-Venus planets, soil moisture also collects at the poles,
both for the 1 bar and the 250 mb atmosphere planets. Likewise, lakes in both
simulations form in both the tropics and polar region, more so over the
highlands for the 1 bar atmosphere but fairly uniformly distributed in
longitude for the 250 mb atmosphere.

##  8 Conclusion

Whether Venus’ original water survived its initial MO stage, or whether
significant water was delivered afterwards, is unknown. It is therefore worth
having a theoretical framework that considers the possibility of an early
habitable Venus as a starting point for designing future observing strategies
that might shed light on Venus’ past. In this spirit, we envision the
following possible climatic evolution for Venus and provide Figure 9 as a
guide:

Figure 9: Graphical representation of Venus’ possible climate history. The
three data points represent the 1 bar atmospheres modeled at those points in
time. The red dashed arrow to the right represents the transition to a
moist/runaway greenhouse and eventually to Venus’ present day surface
temperature and atmospheric density.

1.) Toward the end of the accretion period ($\sim$4.2Ga) Venus would have
cooled rapidly as did Earth, as shown in the work of Valley2002. This would
allow surface water to condense and early oceans to form. Early oceans in turn
could create significant tidal dissipation that would spin down Venus’
rotation rate on a relatively short time scale as shown in work by [Green .
(2019)] and described in Section 5 above. Solid body dissipation may have also
been effective, see Section 5. Slow rotation combined with an early ocean
would provide the necessary ingredients for a dayside cloud deck to emerge,
shielding the planet from high insolation and allowing at least some of the
initial surface water to survive despite the planet being well inside the
conventional inner edge of the habitable zone, as shown in the work of
Yang2014,Way2018.

2.) The carbonate-silicate cycle in concert with interior volatile cycling
would allow CO2 draw-down while N2 was outgassed, eventually reaching a
balance producing an N2-dominated atmosphere with trace amounts of CO2 over
gigayears with pressures ranging from several bars to hundreds of millibars.

3.) We propose that any stable Venusian climate period came to an end at some
period of time (e.g. 0.2Ga to 3 Ga) before the global resurfacing event. As
mentioned previously surface age estimates for Venus range from as young as
$\sim$180 Ma [Bottke . (2016)] to 750Ma [McKinnon . (1997)]. We suggest that
the ignition of multiple large igneous provinces (LIPs) became active around
that time (over a period of 10s or 100s of millions of years). This would not
have been the global LIP as proposed by [López . (1998)] as it is not
necessary and seems to have little support in the community. It is likely that
the eruption of LIPs throughout Earth’s history [Ernst (2014), Ernst . (2019)]
is a random stochastic process. This may imply that multiple _l_ arge scale
LIPs have not occurred simultaneously on Earth by purely random chance, which
is fortuitous for life as we know it today. Venus may not have been as
fortunate. Unfortunately, little is known about Venus’ interior structure
today, much less its initial state and subsequent evolution, so the question
remains as to whether deterministic evolutionary processes in the interiors of
Venus-like planets will inevitably lead to catastrophic changes and
uninhabitable end states.

4.) Multiple large scale LIPs would have warmed the world markedly via the
release of large amounts of CO2 via one or more mechanisms and its greenhouse
effect on the climate [<]e.g.,¿[]Ogden2012 over 10s or 100s of millions of
years. It should be noted that degassing $\sim$ 90 bars worth of CO2 via this
mechanism would be difficult and is an open research topic.

5.) This warming could have hastened a movement toward the onset of a moist
[Kasting . (1993)] and possibly a runaway greenhouse state in which much of
the surface water was lost via photodisassociation accompanied by hydrogen
escape and an oxidized surface. Recall that present day Venus D/H measurements
(if they are accurate) imply .6 to 16% of Earth’s present day surficial water
stores [Donahue . (1997)]. Hence the timescale of total loss would be at least
an order of magnitude faster than that described in Kasting1993b, although
there is still uncertainty about the stratospheric conditions under which a
moist greenhouse occurs [<]e.g.,¿[]Chen2019 and when a runaway greenhouse is
achieved [<]e.g.,¿[]Kasting_Ackerman1986,Goldblatt2013,Ramirez2014. Work by
Grinspoon1993 on the D/H ratio may also support this hypothesis as they state:
“Thus (D/H)obs may be the isotopic signature of a catastrophic resurfacing in
the past 0.5–1 Gyr.”

6.) A major issue with any proposed evolutionary scenario with long lived
surface water reservoirs then arises: what happened to the large quantities of
oxygen expected to be left over from a Venusian ocean? This is the oxygen that
would have remained in the atmosphere after the photo-dissociation of H2O and
the loss of the hydrogen [Watson . (1981)], but not oxygen, via atmospheric
escape. In fact large quantities of abiotically produced O2 (100s to 1000s of
bar) left over in such a scenario has been proposed as an observational
signature in exoplanetary atmospheres on planets that have lost their oceans
[Luger Barnes (2015)]. For our Arid-Venus scenario it may be possible to lose
much of the oxygen via a combination of atmospheric escape [Persson . (2018)]
and absorption by a surface like that of present day Venus
[<]e.g.,¿[]Gilmore2017. However, recently submitted work by Persson2020
demonstrates that 0.3 m of a global equivalent layer of water could have been
lost via atmospheric escape alone in the past $\sim$4Ga. Hence the 0.2 m
global equivalent layer of water in our Arid-Venus scenario fits within this
framework. Yet work by Persson2018,Masunaga2019,Persson2020 and their
estimates for O+ escape rates shows that atmospheric escape alone is not
sufficient to remove larger reservoirs of oxygen left over from the oceans in
our other non-Arid-Venus simulations. It should be noted that the escape
estimates from Persson2020 are distinctly lower than previous work by
Chassefiere1996,Chassefiere1997 where they are mainly concerned with escape
during a more active younger sun. The work of Abe2011 also gives much higher
escape rates estimating that an entire Earth Ocean’s volume could be lost in
600Myr to 14Myr depending on how active the sun is. Some caveats go with this
work in that the Venus atmosphere simulated for our temperate period (more
akin to an Archean Earth atmosphere than modern day Earth or Venus) is N2
dominated with 400ppmv CO2 and 1ppmv CH4 and is very different from that of
modern Venus with thermospheric and exospheric temperatures likely to be
distinct and possibly affecting escape rates [Airapetian . (2017)].

We propose that the large-scale resurfacing evident on Venus today, which took
place over 100s of millions of years, is a possible effective answer. It must
be noted that this would be separate from the earlier (in geological time) LIP
scenario above. Since the Magellan mission it has been known that 80% of the
surface of Venus is relatively young, with estimates ranging from 300-700$Myr$
old as mentioned in Section 4. These newly exposed basalts would be the ideal
sink for large quantities of oxygen (possibly 100s of bars) over 100s of
millions of years. According to [Lécuyer . (2000)] Venus would need to oxydize
a rock layer $\sim$50 km deep to absorb an Earth Ocean’s worth of oxygen, and
they propose a mechanism for doing so while citing the earlier work of
Pieters1986. Note that none of the oceans proposed herein are close to an
Earth’s ocean in volume, hence the number could be much smaller. For example,
work by Grinspoon1993 and Head1992 note that the volume of magma necessary to
cover all pre-existing craters would need to be a global layer $\sim$10 km
deep, and that would be sufficient for the volume of oceans proposed in our
models. Some fraction of the oxygen may actually be deep within Venus’
lithosphere and possibly even within its mantle. This may be consistent with
coronae-related subduction hypotheses [Sandwell Schubert (1992), Davaille .
(2017)] and other ideas about downwelling-associated highlands [<]e.g.¿[see
Fig 2]Head1992, where thicker crustal regions may exhibit orogenesis
[<]e.g.,¿[]Head1990 and sinking of parts of the lithosphere into the mantle
[<]e.g.,¿[]Lenardic1991,Bindschadler1992.

7.) The loss of water would in turn change the planet from an initial
subductive or mobile plate tectonic mode to more of a stagnant lid mode (as on
present day Venus & Mars) since it is currently believed that water plays a
key role in plate tectonics on Earth [<]e.g.,¿[]Grove2012,Lecuyer2014. This
scenario fits in very nicely with the recent work of [Weller Kiefer (2019)]
who give a timescale of order 1Gyr for the transition from a mobile to a
stagnant lid mode on Venus in their simplified model. Without a mechanism to
efficiently cycle volatiles in a stagnant lid mode
[<]e.g.,¿[]Tosi2017,Honing2019, outgassing would have continued without the
major weathering and subduction surface sinks that operate on Earth, hence CO2
and N2 would build up over time to reach the levels we see on Venus today.
Some studies have also shown that even in a stagnant lid mode it is possible
to cycle volatiles, possibly up to gigayears in time
[<]e.g.,¿[]Foley_Smye2018,Godolt2019, but these mechanisms depend on the
initial CO2 budget and the retention of at least some water after cooldown.

8.) This stagnant lid mode may then allow very large mantle upwelling and/or
downwelling centers that would produce some of the features we see on Venus’
surface today produced over hundreds of millions of years, as described most
recently in the works of ¡e.g¿[]Rolf2018,WellerKiefer2019.

Our scenario can also fit within the Fortunian, Guineverian & Atlian periods
proposed in the works of IvanovHead2015b and Airey2017, as it is not possible
to constrain the start of the LIP period we propose with the data we have
today. Our LIP period could easily have concluded in the pre-Fortunian or
Fortunian period $\sim$1.5Ga.

One of the remaining quandaries in our hypothesis is the fact that the 92 bar
atmosphere we see on Venus today must go back at least as far as the age of
the visible surface because there are fewer small craters ($<$35 km in
diameter) to be seen in the Magellan data [<]¿[see Figure 2]Schaber1992.
Certainly smaller craters would be visible if the atmosphere had been
significantly thinner in the lifetime of the observed surface when atmospheric
filtering of smaller impactors would have been less prevalent.

For example, assuming the tesserae are the oldest stratigraphic units why are
there no small craters present if the present day atmosphere is not a
primordial feature from many eons ago? One resolution to this problem could be
that the tesserae are not as old as we think they are, and until we date these
units and the basaltic flats we really won’t know. Secondly, as mentioned
above, there is the possibility that there have been multiple resurfacing
events and the tesserae are left-over from one of the previous events. Neither
is a terribly optimistic scenario if one is hoping that some of tesserae may
be remnant crust from a period of hosting surface water. Finally, a large
impactor may be the cause of the ‘catastrophic’ climate change we propose.
This could have also played a role in resetting the clock on the surface of
Venus reconciling the lack of small craters. In this scenario the LIP
hypothesis plays a partial role in the evolution of Venus’ climate. This would
be similar to what we have seen in the on-going debate regarding the K-Pg
period on Earth [<]e.g.¿[]Hull2020,Schoene2019,Sprain2019. Such an impactor’s
imprint would have long been lost due to Venus’ relatively young surface. Our
comments about impactors in Section 4 apply here as well.

Clearly a great deal more modeling work and more observations are required to
confirm or refute this hypothesis. Did Venus follow the ’canonical’ path with
Earth-like conditions in it’s early history and then experience a moist-
runaway greenhouse due to increasing solar insolation? Did it experience a
longer period of habitability throughout most of its history, with its demise
and present hothouse state the consequence of a series of LIP events releasing
CO2 bound up in the crust as on Earth, and/or released from the deep interior
where CO2 is more easily sequestered [<]e.g.,¿Kuramoto1996? Or did it become
bone dry in an extended magma ocean phase in the first 100$Myr$, as described
in Hamano2013 for Type II planets?

We believe the only scenario we can begin to rule out with the present work is
the ‘canonical path’ since there is no evidence that an early period of
habitability would have been affected by increasing solar luminosity in the
first billion years. In essence, if Venus had habitable surface conditions
with surface liquid water $\sim$4Ga then the same cloud albedo effect that
allows such a scenario would continue for eons. On the other hand we will not
be able to distinguish between the two remaining scenarios until we return to
Venus to make proper noble gas and other elemental and isotopic measurements
at the surface [Baines . (2013)] and better constrain escape processes at the
top of the atmosphere through time. The latter will also rely upon how such
gases escape from present day Earth given the possibility that Venus may have
had a magnetic field in previous epochs, even if it is not clear how important
the magnetic field is to escape processes in general [Gunell . (2018)].
Likewise, whether the actual evolution of the one Venus we can visit is the
ultimate fate of all highly irradiated rocky planets or an accident of an
evolutionary path that might have proceeded differently in other circumstances
[<]e.g.¿[]Lenardic2016 is not known. The stakes are high for answering this
question, since many exoplanets have been discovered in the “Venus zone” just
inside the traditional inner edge of the habitable zones of other stars [Kane
. (2014)]. Efforts to simultaneously characterize the CO2 concentrations and
climates of a number of these exoplanets, combined with a focused
observational strategy for unveiling the history of the “exoplanet next door”
to Earth in our own solar system [Kane . (2019)], will be our best chance to
understand whether the envelope for habitability and the emergence of life is
much broader than usually assumed.

Table 2: Results

IDa | Runtime |  Tempb |  Balancec |  Q${}_{top}^{d}$ |  Q${}_{surf}^{e}$ | Albedo | Albedo | Clouds | Clouds | Clouds
---|---|---|---|---|---|---|---|---|---|---
| years | C |  W m-2 | kg/kg | % | Planetary | Surface | High | Medium | Low
01 | 71 | 262 | 0.01 | 1.36e-05 | 0.08 | 54 | 29 | 54 | 5 | 0
02 | 111 | 151 | 27.21 | 2.67e-05 | 3.74 | 54 | 18 | 33 | 17 | 33
03 | 67 | 121 | 21.31 | 4.50e-04 | 7.87 | 52 | 11 | 45 | 23 | 37
04 | 52 | 120 | 51.28 | 1.74e-04 | 10.01 | 44 | 7 | 36 | 17 | 51
05 | 56 | 123 | 36.27 | 4.28e-04 | 9.16 | 49 | 10 | 46 | 24 | 44
06/01 | 100 | 96 | 0.06 | 1.73e-03 | 1.65 | 51 | 29 | 44 | 2 | 0
07/02 | 100 | 76 | 0.41 | 9.34e-03 | 7.70 | 53 | 16 | 71 | 32 | 21
08/03 | 53 | 61 | 5.92 | 6.49e-03 | 7.50 | 50 | 10 | 72 | 32 | 47
09/04 | 100 | 59 | -0.48 | 6.23e-03 | 6.95 | 52 | 7 | 73 | 35 | 62
10/05 | 60 | 60 | 1.90 | 5.46e-03 | 7.19 | 52 | 9 | 69 | 37 | 53
11/01 | 200 | 73 | 0.05 | 9.03e-04 | 1.70 | 52 | 29 | 59 | 4 | 1
12/02 | 500 | 50 | 0.04 | 5.94e-04 | 2.07 | 55 | 16 | 57 | 10 | 10
13/03 | 200 | 45 | 0.87 | 1.56e-03 | 3.15 | 54 | 10 | 59 | 18 | 43
14/04 | 156 | 53 | 0.11 | 3.47e-03 | 4.81 | 51 | 7 | 66 | 25 | 60
15/05 | 200 | 56 | 0.44 | 4.58e-03 | 5.65 | 52 | 9 | 72 | 28 | 49
16/01 | 300 | 64 | 0.01 | 8.34e-04 | 0.80 | 53 | 28 | 49 | 4 | 1
17/02 | 500 | 44 | 0.06 | 5.68e-04 | 1.40 | 56 | 16 | 48 | 9 | 8
18/03 | 1000 | 37 | 0.13 | 8.18e-04 | 2.07 | 56 | 10 | 55 | 15 | 38
19/04 | 500 | 44 | -0.50 | 1.59e-03 | 3.01 | 54 | 7 | 60 | 17 | 56
20/05 | 500 | 45 | 0.23 | 1.88e-03 | 3.15 | 54 | 9 | 63 | 19 | 44
21/01 | 500 | 22 | 0.02 | 2.52e-05 | 0.36 | 46 | 28 | 45 | 6 | 4
22/02 | 3000 | 21 | -0.04 | 3.08e-05 | 0.61 | 47 | 17 | 50 | 10 | 12
23/03 | 1000 | 10 | -0.05 | 2.53e-06 | 0.70 | 52 | 10 | 56 | 20 | 48
24/04 | 500 | 30 | -0.12 | 3.51e-05 | 2.32 | 45 | 6 | 45 | 14 | 37
25/05 | 1000 | 19 | 0.05 | 5.87e-06 | 1.24 | 48 | 9 | 48 | 22 | 39
26/01 | 300 | 20 | 0.11 | 5.94e-03 | 0.53 | 37 | 29 | 29 | 5 | 9
27/02 | 1000 | 8 | -0.14 | 9.98e-03 | 1.26 | 50 | 20 | 73 | 17 | 27
28/03 | 1000 | 3 | -0.02 | 6.15e-03 | 1.94 | 52 | 13 | 73 | 8 | 13
29/04 | 1000 | 25 | 0.34 | 2.69e-02 | 6.39 | 47 | 6 | 66 | 1 | 0
30/05 | 500 | 10 | 0.08 | 8.58e-03 | 2.76 | 50 | 10 | 71 | 7 | 10
31/01 | 300 | 32 | -0.03 | 1.12e-04 | 0.63 | 52 | 28 | 45 | 7 | 4
32/02 | 297 | 19 | 0.02 | 8.32e-05 | 0.87 | 2 | 1 | 48 | 20 | 26
33/03 | 1000 | 15 | -0.03 | 1.18e-05 | 0.95 | 58 | 10 | 50 | 28 | 47
34/04 | 500 | 36 | -0.03 | 6.41e-04 | 3.32 | 52 | 7 | 38 | 17 | 42
35/05 | 500 | 28 | 0.27 | 1.11e-04 | 2.04 | 55 | 9 | 42 | 26 | 40
36/01 | 50 | 60 | -0.09 | 3.05e-03 | 0.58 | 46 | 29 | 34 | 2 | 1
37/02 | 2000 | 29 | -0.01 | 4.33e-04 | 1.17 | 59 | 17 | 48 | 16 | 23
38/03 | 1000 | 16 | 0.12 | 3.52e-05 | 1.08 | 62 | 10 | 47 | 33 | 47
39/04 | 500 | 38 | -0.07 | 1.37e-03 | 3.57 | 57 | 7 | 42 | 21 | 42
40/05 | 1000 | 34 | 0.35 | 1.10e-03 | 2.92 | 58 | 9 | 46 | 24 | 41
41/01 | 20 | 99 | 0.52 | 3.13e-03 | 0.43 | 32 | 29 | 8 | 0 | 0
42/02 | 8 | 288 | 90.21 | 1.00e-01 | 11.71 | 31 | 17 | 31 | 0 | 0
43/03 | 1000 | 26 | 0.01 | 3.35e-03 | 1.99 | 68 | 11 | 49 | 33 | 54
44/04 | 500 | 38 | 0.23 | 5.52e-03 | 3.75 | 67 | 8 | 51 | 44 | 54
45/05 | 35 | 83 | 32.64 | 1.05e-01 | 30.36 | 66 | 10 | 89 | 19 | 16
D | 2000 | 56 | 0.19 | 5.31e-03 | 8.97 | 44 | 10 | 46 | 39 | 45

aID: The colored numbers correspond to those in Figure 2.
bTemp: Surface temperature in Celsius.
cBalance: These numbers come from 50 year averages unless the Runtime (Column
2) is less than 150 years, then the average is 10 years.
dSpecific humidity in top layer of the atmosphere.
eSurface humidity as percentage of atmosphere.

##  Appendix A Energy balance and temperature

Different simulations reach radiative balance sooner than others, while some
never reach it at all. Herein we plot the energy balance (in units of W m-2)
and surface temperature (in Celsius) as a function of simulation year. This
should allow the reader to have a better grasp of which simulations are
appropriate for a given interest.

Figure 10: Energy Balance (also called Net Radiative Balance) in the left
columns and surface temperature in the right columns as a function of
simulation year for all simulations in this study. Note that not all limits on
the x or y axes are the same.

###### Acknowledgements.

This work was supported by the NASA Astrobiology Program through
collaborations arising from our participation in the Nexus for Exoplanet
System Science (NExSS). Resources supporting this work were provided by the
NASA High-End Computing (HEC) Program through the NASA Center for Climate
Simulation (NCCS) at Goddard Space Flight Center. MJW acknowledges support
from the GSFC Sellers Exoplanet Environments Collaboration (SEEC), which is
funded by the NASA Planetary Science Division’s Internal Scientist Funding
Model. Thanks goes to Rory Barnes for his help with his eqTide program and
Eric Wolf for his discussions involving HITRAN and HITEMP. Thanks also to Wade
Henning,Tony Dobrovolskis, and Kevin Zahnle for their thoughts on spin
evolution. This paper also benefited from useful discussions with Mareike
Godolt, Lena Noack, Richard Ernst, Linda Sohl, Matthew Weller, Christine
Houser, Kostas Tsigaridis, Igor Aleinov, and Glyn Collinson. This paper also
benefited greatly from the detailed comments of three referees: Jim Kasting,
David Grinspoon and Cedric Gillmann. All NetCDF data used in this publication
can be downloaded from the Zenodo open access data portal:
http://doi.org/10.5281/zenodo.3707251

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