Clouds and Hazes of Venus

Abstract

More than three decades have passed since the publication of the last review of the Venus clouds and hazes. The paper published in 1983 in the Venus book summarized the discoveries and findings of the US Pioneer Venus and a series of Soviet Venera spacecraft (Esposito et al. in Venus, p. 484, 1983). Due to the emphasis on in-situ investigations from descent probes, those missions established the basic features of the Venus cloud system, its vertical structure, composition and microphysical properties. Since then, significant progress in understanding of the Venus clouds has been achieved due to exploitation of new observation techniques onboard Galileo and Messenger flyby spacecraft and Venus Express and Akatsuki orbiters. They included detailed investigation of the mesospheric hazes in solar and stellar occultation geometry applied in the broad spectral range from UV to thermal IR. Imaging spectroscopy in the near-IR transparency “windows” on the night side opened a new and very effective way of sounding the deep atmosphere. This technique together with near-simultaneous UV imaging enabled comprehensive study of the cloud morphology from the cloud top to its deep layers. Venus Express operated from April 2006 until December 2014 and provided a continuous data set characterizing Venus clouds and hazes over a time span of almost 14 Venus years thus enabling a detailed study of temporal and spatial variability. The polar orbit of Venus Express allowed complete latitudinal coverage. These studies are being complemented by JAXA Akatsuki orbiter that began observations in May 2016. This paper reviews the current status of our knowledge of the Venus cloud system focusing mainly on the results acquired after the Venera, Pioneer Venus and Vega missions.

Full text

Space Sci Rev (2018) 214:126
https://doi.org/10.1007/s11214-018-0552-z
Clouds and Hazes of Venus
Dmitrij V. Titov1·Nikolay I. Ignatiev2·
Kevin McGouldrick3·Valérie Wilquet4·
Colin F. Wilson5
Received: 30 September 2018 / Accepted: 5 October 2018 / Published online: 27 November 2018
© The Author(s) 2018
Abstract More than three decades have passed since the publication of the last review of
the Venus clouds and hazes. The paper published in 1983 in the Venus book summarized
the discoveries and findings of the US Pioneer Venus and a series of Soviet Venera space-
craft (Esposito et al. in Venus, p. 484, 1983). Due to the emphasis on in-situ investigations
from descent probes, those missions established the basic features of the Venus cloud sys-
tem, its vertical structure, composition and microphysical properties. Since then, significant
progress in understanding of the Venus clouds has been achieved due to exploitation of new
observation techniques onboard Galileo and Messenger flyby spacecraft and Venus Express
and Akatsuki orbiters. They included detailed investigation of the mesospheric hazes in so-
lar and stellar occultation geometry applied in the broad spectral range from UV to thermal
IR. Imaging spectroscopy in the near-IR transparency “windows” on the night side opened
a new and very effective way of sounding the deep atmosphere. This technique together
with near-simultaneous UV imaging enabled comprehensive study of the cloud morphology
from the cloud top to its deep layers. Venus Express operated from April 2006 until Decem-
ber 2014 and provided a continuous data set characterizing Venus clouds and hazes over a
time span of almost 14 Venus years thus enabling a detailed study of temporal and spatial
variability. The polar orbit of Venus Express allowed complete latitudinal coverage. These
studies are being complemented by JAXA Akatsuki orbiter that began observations in May
2016. This paper reviews the current status of our knowledge of the Venus cloud system
focusing mainly on the results acquired after the Venera, Pioneer Venus and Vega missions.
Venus III
Edited by Bruno Bézard, Christopher T. Russell, Takehiko Satoh, Suzanne E. Smrekar and Colin F.
Wilson
BC.F. Wilson
colin.wilson@physics.ox.ac.uk
1ESA/ESTEC, Noordwijk, The Netherlands
2Space Research Institute (IKI), Moscow, Russia
3LASP, University of Colorado, Boulder, USA
4Belgian Institute for Space Aeronomy (IASB-BIRA), Brussels, Belgium
5Dept. of Physics, Oxford University, Oxford, UK
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126 Page 2 of 61 D.V. Titov et al.
Keywords Venus ·Venus atmosphere ·Clouds ·Aerosols
1 Introduction
The last comprehensive review of the clouds and hazes on Venus was published more than
thirty years ago in the Venus book (Esposito et al. 1983). Some updates on the cloud top
structure were made by Esposito et al. (1997) in the Venus-II book. These papers summa-
rized our knowledge about the greatest cloud system on terrestrial planets on the basis of the
results from Venera and Pioneer Venus orbiters and descent probes. The early investigations
provided basic understanding of the Venus cloud system, but lacked details on the cloud
morphology, vertical structure, composition and their variability. The era after the Venera
and Pioneer Venus missions featured remarkable progress in the study of Venus clouds by
the orbiters Venera-15, -16 in 1983, Venus Express in 2006–2014, and Akatsuki since 2016,
balloons and descent probes (Vega-1, -2 in 1985), as well as Galileo (1990) and Messenger
(2007) flyby supported by numerous ground-based observations, numerical modeling, and
theoretical investigations.
This chapter provides an overview of the new results on the Venus clouds and hazes
published after Venus II book. The introduction section gives a brief summary of the mis-
sions, instruments and techniques that enabled progress in cloud investigations after Pioneer
Venus, Venera and Vega missions. Section 2describes vertical structure of the cloud. Sec-
tion 3reviews the cloud morphology through the entire cloud deck. Section 4describes the
progress on observations of microphysical properties of the cloud population. It is followed
by discussion of the modelling of cloud chemistry and microphysics in Sects. 5and 6.Ra-
diative effects directly related to the cloud layer and its variability are described in Sect. 7.
This section complements more detailed description of the radiative energy balance in the
dedicated chapter (Limaye et al. 2018a). Section 8is devoted to lightning in the Venus atmo-
sphere. Sections 9and 10 give a synthesis of the Venus cloud system and outline outstanding
remaining science issues and perspectives for future studies.
The Fourier Spectrometer Experiment (FSE) onboard the Soviet Venera-15, -16 or-
biters (Moroz et al. 1986;Oerteletal.1987) provided about 1500 moderate resolution
(4.5cm
−1and 6.5cm
−1) spectra of Venus thermal radiation in the range of 200–1600 cm−1
(50–6.25 μm) that so far remains unique data set of this kind. The measurements covered
morning and evening sectors mainly in the Northern hemisphere. The data analysis included
retrievals of the temperature structure, calculation of the thermal winds in the mesosphere,
determination of H2OandSO
2abundance at the cloud top, and assessment of the upper
cloud parameters (see Zasova et al. 2007 and references therein).
In 1985 two Soviet Vega spacecraft en route to comet Halley delivered descent probes
and balloons to the Venus atmosphere (Sagdeev et al. 1986; Moroz 1987; Crisp et al. 1990;
Lorenz et al. 2018). The mission science program included in situ measurements of particle
size distribution during descent through the clouds as well as measurements of the backscat-
tering coefficient (nephelometry) on both descent probes and balloons that allowed charac-
terization of aerosols optical properties (see Zasova et al. 1996 and references therein). The
Vega probes also included gas chromatographs, mass spectrometers, and X-ray fluorescent
spectrometers to study chemical composition of the clouds. Some results of those studies
were questionable and controversial (Krasnopolsky 1989).
The NASA Galileo spacecraft en route to Jupiter flew by Venus on February 10, 1990
(Johnson et al. 1991). The NIMS experiment performed pioneering spatially resolved ob-
servations of the Venus night side in the near-IR transparency windows (Carlson et al. 1991)
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Clouds and Hazes of Venus Page 3 of 61 126
discovered by Allen and Crawford (1984). They allowed characterization of spatial variabil-
ity of the total cloud opacity, assessment of microphysical properties in the deep cloud, and
determination of the upper cloud structure. Imaging of Venus at 0.418 μmand0.986 μmby
the SSI camera (Belton et al. 1991) provided the first simultaneous observations of the cloud
morphology and dynamics in the violet and near-IR spectral ranges that sounded the upper
and middle/lower cloud layers respectively.
The ESA Venus Express mission made a breakthrough in our understanding of the clouds
due to its powerful remote sensing payload, long duration of the mission and polar orbit
of the spacecraft (Svedhem et al. 2007,2009,2011; Titov et al. 2006,2009). The highly
elliptical orbit of Venus Express with its apocentre above the Southern pole allowed the
imaging instruments to zoom in on the equatorial and Northern latitudes putting high res-
olution images in global context. Mesoscale images taken by Venus Express from 10000–
15000 km distance covered a significant portion of the planet with spatial resolution com-
parable to the best images from the earlier missions. They provided a close look at the
sub-solar region (i.e. 10–14 hours local solar time). Venus Express was the first Venus or-
biter to have an elliptical orbit with its pericentre at the North pole, providing ideal condi-
tions for planetary scale observations of the Southern high latitudes from apocentre. Venus
Express was also the first orbiter capable of observing the lower clouds of Venus, fully ex-
ploiting near-infrared spectral “windows” on the nightside; these revealed the morphology
and motions of lower/middle clouds, as well as constrained their microphysical properties.
The Visible and Infrared Thermal Imaging Spectrometer (VIRTIS) (Piccioni et al. 2011;
Drossart et al. 2007) performed spectral imaging of the planet in a broad spectral range from
near-UV (0.4μm) to thermal IR (∼5μm), including near-IR spectral transparency “win-
dows” on the night side, thus enabling comprehensive study of the morphology, dynamics,
total aerosol opacity and particle population in the deep cloud. The cloud top altitude and its
variations were retrieved from the VIRTIS observations in the near-IR CO2bands. This was
complemented by assessment of the cloud top structure using the temperature field derived
by the radio occultation experiment VeRa (Häusler et al. 2006). The VeRa sounding also
provided indirect assessment of the cloud base altitude as well as mapped H2SO4vapour
abundance in the deep cloud (Oschlisniok et al. 2012).
The Venus Monitoring Camera (VMC) provided continuous monitoring of the cloud
layer in four narrow-band filters from UV to near-IR with spatial resolution ranging from
50 km to few hundreds of meters (Markiewicz et al. 2007,2011) resulting in detailed charac-
terization of the cloud top morphology and wind pattern as well as optical and microphysical
properties of the upper cloud.
The vertical structure and microphysical properties of the upper haze were studied in
detail by the SPICAV-SOIR spectrometer onboard Venus Express (Bertaux et al. 2007). The
experiment performed medium to high spectral resolution spectroscopy in the UV (118–
320 nm), visible and near IR (0.7–1.7μm) and IR (2.2–4.3μm) ranges in stellar and solar
occultation geometry. The observations covered altitudes from 160 km to about 70 km that
includes mesosphere and lower thermosphere. Due to large air mass factor the measure-
ments were very sensitive to aerosol properties at these altitudes. The observations provided
complete latitudinal coverage and characterized spatial and temporal variability of the upper
haze.
The task of monitoring of the Venus atmosphere has been taken over from Venus Express
by the JAXA Akatsuki spacecraft (Nakamura et al. 2007) inserted in orbit in December
2015. The orbiter carries a set of cameras imaging the planet in the broad spectral range from
UV to thermal-IR to study cloud morphology and atmospheric dynamics at different levels
within the cloud deck. The UVI camera (Yamazaki et al. 2018) and the IR1 (Iwagami et al.
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126 Page 4 of 61 D.V. Titov et al.
2011)andIR2(Satohetal.2017) imagers continue investigations of the aerosol population
at the cloud top and in the mesosphere. The thermal-IR camera (LIR) characterizes global
cloud top temperature distribution (Fukuhara et al. 2017).
While many of the findings are still in development and so are too incomplete to include
in this review, we do note that the early results from the mission have revealed new insights
into the atmosphere of Venus. For example, the nearly planetary scale bow wave seen in the
cloud tops by LIR suggests a direct coupling between the flow near the surface and that at
70 km altitude (Fukuhara et al. 2017). In addition, measurement of a persistent but inter-
mittent zonal jet at 50 km altitude observed in IR2 data suggests that the Venus atmosphere
exhibits changes of a magnitude, and on spatial and temporal scales that we would term
as “weather” on Earth (Horinouchi et al. 2017). These findings and the promise of more to
come show that there is still very much to learn about the workings of the Venus atmosphere
and clouds.
Ground-based observations of Venus appeared to be an effective tool supplementing or-
biter studies. Several ground-based observation campaigns were carried out in coordination
with Venus Express investigations (see e.g. Lellouch and Witasse 2008). Tavenner et al.
(2008) presented a multi-year sequence of nightside near-IR (2.26 μm) images and derived
global mean cloud pattern. Arney et al. (2014) provided spatially resolved measurements of
minor species, cloud opacity and acid concentration in the near-IR “windows” on the Venus
night side. Encrenaz et al. (2012) reported thermal mapping of Venus with high spectral
resolution to determine spatial and temporal variations of H2OandSO
2at the cloud tops,
while Marcq et al. (2006) and Chamberlain et al. (2013) mapped water abundance in the
lower atmosphere.
2 Vertical Structure of the Cloud
Background Venus is completely shrouded by the clouds that form the largest aerosol
system among the terrestrial planets. The basic knowledge about its vertical structure was
established by Venera, Pioneer Venus, and Vega descent probes, the last of which conducted
its mission in 1985 (Esposito et al. 1983; Ragent et al. 1985). There have been thirteen
instrumented entry probes which have returned information about the vertical structure of
the Venus clouds. There are some characteristics common to all the profiles, but there are
also some significant differences. In particular, the Pioneer Venus Large (Sounder) probe
(Knollenberg and Hunten 1980) and the Vega-1, -2 LSA aerosol particle size spectrometer
(Zhulanov et al. 1986), all show a sharp increase in the cloud extinction just below 50 km
suggesting an optically thick layer at 47–50 km. Some descent probes did not see this fea-
ture, suggesting strong variability of the deep cloud structure. There is also considerable
evidence for hazes below the main cloud layer. Both the Pioneer Venus Large Probe and
Vega ISAV UV spectrometer showed discrete cloud layers near 46 km altitude, distinct from
the lower cloud. The ISAV spectrometer found a second such layer at 43 km altitude (Knol-
lenberg and Hunten 1980; Bertaux et al. 1996). Pioneer Venus and Vega probes also show
evidence for further hazes extending down to 30 km altitude (Knollenberg and Hunten 1980;
Moshkin et al. 1986; Gnedykh et al. 1987). These layers all exist at an altitude below that
where one would expect to find liquid sulphuric acid at equilibrium. Finally, a reanalysis of
the Venera-13, -14 spectrophotometer data by Grieger et al. (2004) tentatively suggested an
evidence for a detached aerosol layer with extinction of ∼1.5km
−1at visible wavelengths
1–2 km above the surface.
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Clouds and Hazes of Venus Page 5 of 61 126
Tab le 1 Parameters of the Venus cloud system (from Esposito et al. 1983). The modes of aerosol population
are based on the data from Particle Size Spectrometer experiment on the PV Large Probe (Knollenberg and
Hunten 1980). Note that this terminology should be used with caution in representing other data sets
Region Altitude range (km) Optical depth, τ
(at 0.63 μm)
Mean diameter (μm) Average number density
(Ncm−3)
Upper haze 70–90 0.2–1.0 0.4 500
Upper cloud 56.5–70 6.0–8.0 Mode 1: 0.4 1500
Mode 2: 2.0 50
Middle cloud 50.5–56.5 8.0–10.0 Mode 1: 0.3 300
Mode 2: 2.5 50
Mode 3: 7.0 10
Lower cloud 47.5–50.5 6.0–12.0 Mode 1: 0.4 1200
Mode 2: 2.0 50
Mode 3: 8.0 50
The atmospheric structure derived from the earlier missions is summarised in Fig. 1.This
structure is typical for the low latitudes. The main cloud deck extends from about 48 km up
to ∼70 km. It can be subdivided in three layers according to the behaviour of extinction
coefficient and particle population. The upper cloud (57–70 km) is populated by submicron
(r1∼0.2μm) and micron size (r2∼1μm) particles (Table 1), called mode 1 and 2 respec-
tively. This is the altitude range where the photochemical “factory” producing sulphuric acid
from SO2and H2O is located. Another species present solely in the upper cloud is the mys-
terious UV-blue absorbers whose inhomogeneous vertical and spatial distribution creates
well-known markings on the cloud top that are routinely used to study the cloud morphol-
ogy and atmospheric dynamics (see Sect. 3). This species strongly absorbs at 0.3–0.5μm
and is responsible for absorption of about half of the solar energy the planet receives from
the Sun. The upper cloud is also the region of the strongest variability of the temperature
structure. The tropopause is located approximately at the base of this region at about 60 km
(Fig. 1) (Limaye et al. 2018a). Zonal wind speed reaches its maximum of 100–120 m/sin
the upper cloud (Fig. 1) indicating that the cloud strongly affects atmospheric circulation
(Sánchez-Lavega et al. 2017).
The upper and middle clouds are often separated by a 1–2 km gap with reduced extinc-
tion located at ∼56 km. Below that level the cloud density gradually increases with depth
reaching its maximum at ∼50 km. The separation between the middle and the lower clouds
is not very clear. This region is characterized by tri-modal particle distribution with typi-
cal radii of 0.15–0.2μm (mode 1), 1–1.25 μm (mode 2) and 3.5–4.0μm (mode 3). The
existence of a separate mode 3 of large particles is still a controversy: the large particles
might be just a “tail” of mode 2 rather than a separate mode (Toon et al. 1984). As in the
upper layer sulphuric acid was found to be the major aerosol constituent in the middle and
lower clouds, although significant elemental abundances of chlorine and phosphorous were
also found at these altitudes (Andreichikov 1987) (Sect. 5). The temperature gradient here
is close to adiabatic lapse rate and the stability parameter (difference between the mea-
sured temperature gradient and the adiabatic lapse rate) is close to zero (Seiff et al. 1985;
Limaye et al. 2018a) suggesting that convection dominates the energy and material transport
in this part of the cloud, in contrast to the convectively stable upper layer.
Extended layers of fine aerosols are observed both above and below the main cloud deck
(Fig. 1). The upper haze fills the mesosphere up to ∼100 km altitude with evidences of
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126 Page 6 of 61 D.V. Titov et al.
Fig. 1 Vertical structure of the Venus clouds as derived by the Venera and Pioneer Venus descent probes.
The color lines show mean temperature profiles at low (red), middle (blue) and high (green) latitudes with a
single black line representing the temperature structure below 30 km. The typical vertical profile of aerosol
extinction (Ragent et al. 1985) is shown on the left. The static stability profile from the VIRA model (Seiff
et al. 1985) is shown in the middle
detached layers (Fig. 2). The haze is presumably composed of sulphuric acid. The lower
haze extends down to ∼33 km, far below the level of sulphuric acid thermal decomposition.
Descent probes also provided some evidence for thin aerosol layers near the surface.
Later observations will be reviewed here in turn.
Upper Haze (70–90 km) FSE experiment onboard Venera-15, VIRTIS, VeRa, VMC and
SPICAV-SOIR experiments onboard Venus Express, and Akatsuki cameras provided a more
detailed view on the vertical structure of the Venus cloud system, in particular on its spatial
and temporal variability.
Haze is ubiquitous in the mesosphere and extends from the cloud top (∼70 km) up to
∼110 km. Sulphuric acid particles make up most of the upper haze. The upper haze prop-
erties were probed through polarimetry (Kawabata et al. 1980; Esposito and Travis 1982;
Braak et al. 2002) and dayside limb scans (Lane and Opstbaum 1983; Krasnopolsky 1983),
all at relatively short wavelengths. Venus Express and Akatsuki missions provided signifi-
cant progress in remote studies of the Venus upper haze and its variability in a wide spectral
range. Solar and stellar occultation as well as observations in limb geometry enable detailed
characterization of aerosol vertical distribution. The near-IR spectrometer SPICAV-SOIR
sounded the Venus mesosphere between 0.22 μmand4μm in solar occultation geometry
suggesting gradual increase of slant opacity and local extinction coefficient (β,km
−1) with
decreasing altitude (Wilquet et al. 2009,2012; Luginin et al. 2016). SPICAV-UV performed
stellar occultations at 0.22–0.3μm on the night side (Wilquet et al. 2009). VIRTIS imag-
ing spectrometer measured limb spectra in the IR range (1.05–5.19 μm). Thermal radiation
from the cloud top scattered by the mesospheric haze increased by a factor of ∼10 between
90 and 82.5 km (de Kok et al. 2011). The infrared cameras onboard Akatsuki are delivering
images of the full Venus disk (Satoh et al. 2015,2017).
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Clouds and Hazes of Venus Page 7 of 61 126
Fig. 2 Typical profiles of
aerosol extinction derived from
SPICAV-SOIR observation on 19
August 2007 (6 am, 70 N)
onboard Venus Express at
different wavelengths (see colour
legend) suggesting detached
aerosol layer at 78–84 km
Figure 2shows examples of aerosol extinction vertical profiles derived from the solar
occultation sounding onboard Venus Express. Fine upper haze extends throughout the meso-
sphere up to ∼100 km with extinction coefficient decreasing by more than two orders of
magnitude over 25 km altitude range. In some cases, observations with vertical resolution
better than 2 km when the spacecraft was close to the planet allowed identification of de-
tached aerosol layers (Wilquet et al. 2009). Optically thin haze layers were observed at
80–85 km in about 60% of high-resolution profiles (Luginin et al. 2016).
Figure 3compares the aerosol scale height Haderived from different experiments. The
upper haze region (75–90 km) is rather uniform over the planet and is characterized by
Ha=2.8±1.2 km for 20–80◦N(deKoketal.2011; Luginin et al. 2016), with no evidence
of morning-to-evening variability. At high polar latitudes (>80◦N), the scale height was
found to be larger (Ha=4.4±1.0 km). Below 75 km at the cloud top the scale height
increases to 4–5 km at low and middle latitudes, while in the “cold collar” and some polar
regions the scale height can decrease to below 2 km indicating sharp upper boundary of the
cloud. The sharp cloud top is usually associated with the regions of strong thermal inversions
(Limaye et al. 2018a) suggesting physical relation between these two features (see Sect. 7).
Comparison to the gaseous scale height Hg∼4.5 km shows that at least in low latitudes
aerosol at the cloud top is well mixed with the gas. Decrease of the scale height of the upper
haze above 75 km indicates that aerosol production occurs at the cloud top.
Understanding of the upper haze variability is of great importance for chemistry and ra-
diative balance of the mesosphere. Early studies of spatial and temporal variations of Venus
polarization by the Pioneer Venus Orbiter (PVO) between 1978 and 1990, revealed latitudi-
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126 Page 8 of 61 D.V. Titov et al.
Fig. 3 Scale height of the mesospheric haze at low latitudes derived from observations: 1 (black)—
SPICAV-IR: solid/dashed correspond to latitude ranges 60◦–80◦N/>80◦N (Luginin et al. 2016,2018),
2 (green)—VIRTIS/VEX (de Kok et al. 2011), 3 (orange)—NIMS/Galileo (Roos-Serote et al. 1993), 4 (yel-
low)—ground-based observations (Roos-Serote et al. 1996), 5 (cyan)—VIRTIS&VeRa/VEX: solid/dashed
correspond to low-middle/high latitudes (Lee et al. 2012), 6 (red)—FSE/Venera-15 (Zasova et al. 1993). Blue
line shows gaseous scale height according to the VIRA model (Seiff et al. 1985)
nal variations of an order of magnitude in haze opacity indicating that the haze is likely of
photochemical origin. The observations also suggested temporal variations on the order of
hundreds days and a long-term declining trend of the haze opacity over 11 years of the PVO
mission (Kawabata et al. 1980; Braak et al. 2002).
SOIR and SPICAV-IR observations of aerosol extinction over 8 years of the Venus Ex-
press mission established climatology of the mesospheric haze (Fig. 4) (Wilquet et al. 2012;
Luginin et al. 2016). The analysis confirmed latitudinal dependence of extinction, the extinc-
tion coefficient being at least one order of magnitude greater at the equator than that at the
poles. At 0–70◦latitude, the level of slant opacity τ∼1at3μm is located 5–10 km lower on
the morning side as compared to the evening side of the terminator. Thus the mesospheric
haze demonstrates the latitudinal trend similar to that of the cloud top (see Figs. 7,8). In
the polar regions the morning-evening difference is negligibly small. Latitudinal variations
of the altitude of τ∼1 level between polar and equatorial latitudes reach about 12 km in
the morning and almost 20 km in the evening. All this is consistent with photochemical
production of the mesospheric haze on the day side.
Both short and long-term variability of the upper haze was revealed during the eight years
of the Venus Express mission. The extinction variations can reach an order of magnitude on
the time scale of a few Earth days (Wilquet et al. 2012). Observations of temporal variations
of the Venus polarization during the Pioneer Venus Orbiter mission were revisited by Braak
et al. (2002). The haze particle column density was confirmed to decrease gradually by a
factor of ∼5 during 12 years of the mission.
The haze extinction coefficient at low latitudes (40◦S–40◦N) increased by more than one
order of magnitude during the first 1000 orbits of the Venus Express (Fig. 5). The upper
haze did not show any systematic or periodic variations between 2006 and 2014 except for
strong increase of low latitude haze in the beginning of the mission. SPICAV-UV and SOIR
observations suggested that SO2mixing ratio at the cloud top decreased from equator to
pole (Belyaev et al. 2012; Marcq et al. 2011). SO2abundance increased in the beginning of
the mission and then dropped. These temporal variations might be due to a multiple volcanic
eruptions or to changes in the atmospheric circulation (Marcq et al. 2013). As a result, the
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Clouds and Hazes of Venus Page 9 of 61 126
Fig. 4 Altitude of the reference slant optical depth (τ=1at3μm) as a function of absolute latitude for solar
occultation observations at 6 am (blue) and 6 pm (red) (updated from Wilquet et al. 2012)
SO2concentration at cloud-top and the mesospheric haze opacity in the upper haze only
partly correlate.
Upper Cloud (56.5–70 km) The cloud top can be observed from space in the broad spec-
tral range from UV to thermal IR. The upper cloud boundary is rather diffuse in low and
middle latitudes and becomes considerably sharper at high latitudes. The cloud top alti-
tude varies with wavelength due to changing extinction properties of aerosols. Ragent et al.
(1985) summarized that the altitude of the upper cloud boundary is located at 65 to 70 km
and varies from equator to pole. It is depressed by about 3 to 4 kilometers near the “cold
collar” from its lower latitude values, and then rises by 2 to 3 kilometers higher at about
80◦N, forming a lip into the polar region.
Venus Express observations allowed detailed characterization of the cloud top altitude.
Spectroscopic observations by VIRTIS and SPICAV provided a reliable tool to monitor
the location of the upper cloud boundary and its variability. Depth of the carbon dioxide
absorption bands in the infrared range is proportional to the total number of CO2molecules
on the line of sight and, thus, depends on effective path of radiation in the atmosphere. This
path is a function of the cloud top altitude, its vertical structure, aerosol optical properties,
atmospheric temperature and pressure, and geometry of observations. The cloud top altitude
is defined as the altitude of the unit optical depth (τ=1 level) and therefore is wavelength
dependent (Fig. 6). However, in a wide spectral range from UV to 1.6μm the dependence is
rather weak for 1 μm sized particles that form the main part of the cloud particle population
in the upper cloud in low and middle latitudes.
Ignatiev et al. (2009)used1.6μmCO
2band in the VIRTIS-M spectra to map the cloud
top altitude from all available dayside observations. Later Cottini et al. (2012,2015)used
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126 Page 10 of 61 D.V. Titov et al.
Fig. 5 Long-term variations of the upper haze volume extinction coefficient over the course of Venus Express
mission for two different latitude bins. Mean of the local extinction at 3 μm and altitude of 80 km for each
season of solar occultation for 0◦–40◦N (left) and 80◦–90◦N (right) bin of absolute latitude (updated from
Wilquet et al. 2012)
Fig. 6 Wavelength dependence
of the cloud top altitude (i.e.
τ=1 level) calculated for several
Venus cloud models (Ignatiev
et al. 2009)
2.5μmCO
2band in the spectra measured by the high resolution channel of VIRTIS to de-
rive the cloud top altitude. These results suggested that the cloud top altitude reported by
Ignatiev et al. (2009) were affected by a systematic error and should be corrected by approxi-
mately −2 km. SPICAV measurements at 1.48 μm also supported this suggestion (Fedorova
et al. 2016). Summarizing these three sets of measurements we come to the following con-
clusions. In low and middle latitudes the cloud top at 1.5μm is located at 72 ±1km.It
decreases poleward of ±50◦and reaches 61–67 km in the polar regions (Fig. 7). No consid-
erable local time variations were observed. The average latitudinal profile of the cloud top
altitude is smooth, although instantaneous profiles have local maxima of several hundred
meters over the average trend at 50◦–70◦latitudes with a typical size of 10◦along meridian
(Ignatiev et al. 2009; Cottini et al. 2012,2015). Fast variations at the scale of about 1 km
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Clouds and Hazes of Venus Page 11 of 61 126
Fig. 7 Mean cloud top altitude
as a function of latitude and local
time. Linear patterns are the
traces of VIRTIS-M image
frames (from Ignatiev et al. 2009
with systematic error corrected)
occur in tens of hours, while larger long-term variations of about several kilometers have
been observed only at high latitudes. In low latitudes the cloud top altitude averaged over
hundred day periods is remarkably stable.
Cottini et al. (2015) did not find any systematic correlations between the cloud top alti-
tude, water vapor abundance and brightness at 0.375–0.385 μm. However, dark UV features,
with characteristic size of a few degrees of latitude (i.e. several hundred kilometers), tend
often to coincide with enhanced cloud density (or, equivalently, with higher cloud tops) and
bright features with less dense (or deeper) clouds. This result seems to contradict the gen-
eral understanding derived from global Pioneer-Venus and Venus Express imaging that the
UV-dark material is located deeper in the upper cloud and is being brought to the cloud top
by dynamical mixing (Esposito et al. 1983; Titov et al. 2008) although this conclusion was
related to the global UV pattern and may not take into account local processes.
At longer wavelengths in the thermal IR range the cloud top (τ=1) level is located sev-
eral kilometers deeper than in the near IR region, but demonstrates similar latitudinal trend
(Zasova et al. 2007; Lee et al. 2012;Hausetal.2013;Hausetal.2014)(Fig.8). The aver-
age latitudinal profile of the cloud top altitude is symmetric with respect to equator. In both
hemispheres it starts decreasing at ∼30◦.At50
◦–60◦its latitudinal gradient becomes larger
and the cloud top altitude reaches its minimum at the pole. In the far IR range (∼30 μm) the
sulfuric acid absorption is much smaller than that at shorter wavelengths, and so the cloud
top altitude is located considerably deeper (e.g. by about 10 km at low latitudes) than at
shorter wavelengths (Fig. 8). Interestingly the cloud top strongly descends from equator to
pole in the wavelengths range 1–8 μm that sounds the upper cloud, while there is almost
no decrease in cloud top altitude at 30 μm (open circles in Fig. 8) that probes the middle
cloud. This suggests that the upper cloud shrinks in vertical direction towards the pole while
the middle cloud does not change its structure. The second peculiarity seen in Fig. 8is that
poleward from the “cold collar” the 8 μm cloud top (filled circles) is located deeper than that
at 5 μm (diamonds) while the trend is opposite in low latitudes. Lee et al. (2012) argued that
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126 Page 12 of 61 D.V. Titov et al.
Fig. 8 Latitude dependence of the cloud top altitude overplotted on the latitude–altitude temperature field
derived from analysis of VIRTIS-M dataset (Haus et al. 2014): dashed line—1.5μmCO
2band (Ignatiev
et al. 2009) with systematic shift corrected (Cottini et al. 2012); squares—2.5μmCO
2band (Cottini et al.
2015); filled and empty triangles—1 μmand5μm, respectively, derived for the northern hemisphere and
mapped to the southern hemisphere symmetrically with respect to equator (Haus et al. 2013); filled and
empty circles—8.2 and 27.4μm, respectively (Zasova et al. 2007)
this behaviour might indicate larger particle sizes in the polar regions than at low latitudes,
since large particles (r∼3–4 μm) have smaller extinction efficiency and thus deeper cloud
top at 8 μmthanat5μm. A similar trend was found by Garate-Lopez et al. (2015).
Simultaneous observations by the VIRTIS and VMC instruments onboard Venus Express
provide an opportunity to correlate the cloud top altitude pattern with the UV markings at
global scale (Fig. 9). The UV bright and dark mesoscale features can be traced also in
the cloud top altimetry maps as well as in the cloud top temperature fields. The UV dark
spiral and circular features usually present at −70◦are clearly seen in the cloud altimetry
maps as variations of several hundred meters overlaid on the global descent to the pole
(Fig. 9). Contrary to the analysis of the Pioneer Venus polarization measurements (Esposito
and Travis 1982) and conclusions derived by Titov et al. (2008), the dark polar features are
located higher or often correspond to the increasing latitudinal gradient of the cloud top
altitude. The centre of the polar depression in the cloud top altitude always coincides with
the “eye” of the polar vortex observed by VIRTIS at thermal IR wavelengths (Figs. 9,10)
(Piccioni et al. 2007). The “eye” that usually appears almost featureless at UV wavelengths
has complex structure in the cloud altimetry map that perfectly correlates with the cloud
top temperature. “Hot” spiral arms that have almost the same temperature as the core are
located higher and characterized by strong gradient of the cloud top altitude, thereby being
the boundary of the polar vortex. This feature will be discussed in more detail in Sect. 3.
The sharpness of the cloud top boundary, characterized by aerosol scale height, is latitude
dependent. In low latitudes the upper cloud is rather diffuse, while in the “cold collar” and
polar regions the cloud top boundary can be very sharp. As well as the cloud top altitude,
the scale height is wavelength dependent. However, although the cloud top altitude in the
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Clouds and Hazes of Venus Page 13 of 61 126
Fig. 9 VMC UV images with overplotted cloud altimetry maps. Note correlation of the cloud altimetry and
UV dark and bright features (Ignatiev et al. 2009)
UV and thermal IR differ by 15 km, the cloud scale height evaluated from Pioneer Venus
(Koukouli et al. 2005), Venera-15 (Zasova et al. 1993; Koukouli et al. 2005), Galileo (Roos-
Serote et al. 1993) and Venus Express (Lee et al. 2012) observations in nadir and solar
occultation geometry are in good agreement and indicate the trend for the aerosol scale
height to decrease with altitude in the mesosphere (Fig. 3). The observations also show a
remarkable latitudinal trend (Fig. 11). The cloud top scale height strongly decreases from
∼4 km (that is similar to the gaseous scale height) at low-to-middle latitudes to ≤1km
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126 Page 14 of 61 D.V. Titov et al.
Fig. 10 Correlation of the fine
structure of the vortex eye
(orange image, 5 μm) with the
cloud top pattern (blue isolines).
Altimetry data are absent in the
nightside (upper part) and in a
bright region of the dayside
(lower part), where the signal is
saturated. Contours are drawn
every 0.2 km starting from
67.8 km in the center (Ignatiev
et al. 2009)
Fig. 11 Latitude dependence of
the cloud scale height from
Zasova et al. (1993)andLee
et al. (2012)
in the “cold collar”. Further poleward the aerosol scale height varies from 1–1.5 km in the
“hot dipole” and polar regions to >4 km in the transition region (Fig. 11) suggesting strong
spatial variability of the cloud top structure at high latitudes.
Middle and Lower Clouds (47.5–56.5km) The Soviet Vega missions in 1985 had been
the last so far to conduct in situ investigations of Venus by descent probes and balloons
(Sagdeev et al. 1986). The measurements of aerosol properties by ISAV-A particle size
spectrometer and nephelometer (Moshkin et al. 1986; Gnedykh et al. 1987) and LSA photo-
electric aerosol counter (Zhulanov et al. 1986) on Vega descent probes in general confirmed
the earlier results (Knollenberg and Hunten 1980), but also revealed substantial differences.
Gnedykh et al. (1987) pointed out that above 55 km the ISAV-A measurements could have
been compromised by smoke particles originated from the pyro-jettisoning of the probe’s
heat shield. Below this level, i.e. in the middle and lower cloud, the particle size distribution
was characterized by two modes. Mode 1 can be approximated by the power law (n∼1/rα)
in the particle radius range 0.25–2.5μm with an exponent of 5 ±1forVega-1and4±0.5for
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Clouds and Hazes of Venus Page 15 of 61 126
Vega-2. Mode 2 was composed of the particles with r=1–2.5μm. The particles with larger
radii were rarely seen, and no separate mode of larger particles was detected. Both ISAV-A
and LSA experiments detected much smaller number density in the middle and lower cloud
layer as compared to Pioneer Venus. The mode 2 particles were about an order of magni-
tude (N<10 cm−3) less numerous. They were found to be spherical with refractive index
of 1.4±0.05. Two groups of particles were distinguished within mode 1: about 80% of the
population had refraction index m=1.4±0.1 while about 20% had much higher value of
refractive index (m=1.7±0.1). The measurements also suggested non-sphericity of small
particles.
Gnedykh et al. (1987) used the aerosol properties derived from the particle size spec-
trometer to model the backscattering nephelometer signal and vertical profile of thermal
radiation leaking from the surface—the quantities measured by the same instrument. The
analysis suggested a presence of large number of small particles with size below ISAV-A
detection limit (r<0.25 μm). The results indicated that the haze with high refractive index
m=2 and number density of 5 ·104–5 ·105cm−3extends down to ∼35 km. The esti-
mated mass density of the haze was in the range of 0.1–2 mg/m3. The authors mentioned
that similarly dense lower haze was observed during Venera-8 descent in 1972. Since Vega
and Venera-8 probes landed on the night side that might indicate development of dense sub-
cloud haze at night. On the other hand, the nephelometric profiles recorder by two PV night
probes did not support this finding. We note that analysis of VMC/Venus Express observa-
tions of glory also led Petrova et al. (2015) and Shalygina et al. (2015) to the conclusion
about presence of the particles with high refractive index at the cloud top.
Aerosol properties in the middle cloud were sounded in situ by the nephelometer experi-
ment onboard Vega-1 balloon that floated at 53.5–55 km altitude for about 46 hours having
covered the distance of ∼11000 km (about 105◦longitude) driven by zonal winds (Sagdeev
et al. 1986;Lorenzetal.2018). For most of the flight duration Sagdeev et al. (1986), Ragent
et al. (1987) and Crisp et al. (1990) reported unbroken clouds with scattering properties in
agreement with the other descent probes. Over the range of altitudes traversed by the Vega-1
balloon during the course of the flight the measured backscattering coefficient varied from
0.8·10−4m−1sr−1to 1.8·10−4m−1sr−1with general trend to decrease with altitude sug-
gesting the particle scale height of about 3 km. In some regions a lack of anti-correlation with
altitude indicated small scale variability in the cloud structure. During the period of greatest
convective activity the measurements suggested about a factor of two greater backscatter-
ing. Ragent et al. (1987) tentatively attributed these events to moderate increase of number
density of large particles admixed from lower regions by convective motions. During pe-
riods of minor convective activity the observations indicated much small-scale backscatter
fluctuations with time scale of about 5 min.
While Venus Express has brought a wealth of information on the upper hazes and cloud
top structure, it has provided only indirect constraints on the vertical structure of the deep
cloud from the measurements of the wings of the nightside infrared spectral “windows”. As-
suming sulphuric acid composition of the clouds and particles size distribution from Knol-
lenberg and Hunten (1980), Barstow et al. (2011) found that the observed spectral variations
of the nightside emissions can be explained by changes in the altitude of the cloud base
from 46 km at 50◦Sto42kmat75
◦S. However, these results should be taken with caution
since the nightside emission is sensitive to several parameters of the cloud and distinguish-
ing between them can be ambiguous. For instance, Haus et al. (2013) did not report a need
for variable cloud base to fit VIRTIS/Venus Express spectra. Little variation in the cloud
structure was observed as a function of local solar time and longitude. The total opacity
of the clouds was derived from the radiance measured in the near-IR transparency “win-
dows”. The opacity at 1 μm averaged over the globe is 34.7 (Haus et al. 2013). According to
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126 Page 16 of 61 D.V. Titov et al.
NIMS/Galileo observations the total opacity at 1.7 and 2.3μmrangesfrom25to40(Grin-
spoon et al. 1993). Satoh et al. (2009) used a small subset of VIRTIS/Venus Express data in
the 1.74 μm “window” to assess properties of the lower haze at 30–40 km. They tentatively
suggested that the lower haze has total opacity of 0.5–3 and consists of very small particles.
The total aerosol opacity is 30–50.
Further indirect evidence that the cloud base descends by several kilometers towards the
pole can be found in VeRa/Venus Express measurements of static stability (Tellmann et al.
2009;Limayeetal.2018a). This parameter is quite sensitive to the presence of infrared
opacity sources such as clouds. Zero stability is likely to indicate presence of dense clouds.
VeRa radio occultation sounding suggested zero stability atmosphere extending down to
50 ±2 km at low and middle latitudes and about 5 km deeper at polar latitudes (>70◦)
thus indicating presence of dense clouds down to ∼45 km in the polar regions. Similarly,
mapping the nightside near-IR emissions also indicated greater total cloud opacity in the
polar region (see Fig. 20 in Sect. 3). More direct constraints on the cloud base altitude may
eventually be obtained from the measurements of sulphuric acid vapour abundance at 40–
60 km altitude by the radio occultation experiment VeRa/Venus Express (Oschlisniok et al.
2012).
3 Cloud Morphology in 3-D
Since extinction of the Venus atmosphere is strongly spectrally dependent, imaging of the
planet at different wavelengths is a powerful tool to sound different altitudes and proper-
ties of the cloud layer. At visible wavelengths Venus appears as a bright featureless white
disc due to high albedo uniformly across visible wavelength (Fig. 12, upper left). The cloud
morphology is well pronounced if observed at a specific wavelength of the unknown UV
absorber (365 nm) of which inhomogeneous distribution at the cloud top (∼60 km) is re-
sponsible for the pattern (Fig. 12, upper right). Thermal IR imaging gives access to the cloud
top temperature and its variations that can reach 40–50 K (Fig. 12, lower right). Imaging in
the narrow spectral transparency “windows” in the near-infrared range on the night side pro-
vides back illuminated view of the clouds (Fig. 12, lower left). This adds vertical dimension
to the cloud morphology, since the patchy pattern observed at these wavelengths is created
by opacity inhomogeneities in the deep cloud (∼50 km).
Venus Express observations from UV (Markiewicz et al. 2007,2011) to thermal IR
(Drossart et al. 2007; Piccioni et al. 2011) provided remarkable progress in understand-
ing of the cloud morphology. The broad spectral coverage enabled sounding of the Venusian
clouds from their tops to deep layers. The images revealed a great variety of features at dif-
ferent spatial and temporal scales, strong variability of the cloud patterns and unexpected
correlations between the features seen at different altitudes. They also provided excellent
material to characterize general circulation and dynamical properties of the Venus atmo-
sphere within the cloud deck (Sánchez-Lavega et al. 2017). This section describes in detail
the cloud morphology observed by Venus Express and reveals morphological relations be-
tween the images observed at different wavelengths.
Near-UV brightness contrasts that reach 20–30% are produced by inhomogeneous distri-
bution of an unknown absorber mixed within the sulphuric acid aerosol in the upper cloud
layer. Radiative transfer modelling by Haus et al. (2016) suggests that apparent changes in
the near-UV albedo of ∼10% require more that 25% changes in total abundance of the
absorber. The UV markings have been routinely used to study the cloud top morphology.
Rossow et al. (1980) described basic types of the cloud features seen in the Pioneer Venus
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Clouds and Hazes of Venus Page 17 of 61 126
Fig. 12 Examples of Venus
views at different wavelength:
equatorial view from
MESSENGER flyby (upper left);
mosaic composed of UV
(365 nm) (gray) and near-IR
(1.7μm) (red-black) images
simultaneously taken by Venus
Express in the middle and high
latitudes of the Southern
hemisphere (upper right); false
colour equatorial view in the
2.3μm spectral transparency
“window” by Galileo (lower
left); false colour thermal IR
(11.5μm) image of the Southern
hemisphere by Pioneer Venus
(lower right). (Credits NASA and
ESA)
Orbiter Cloud Photopolarimeter (OCPP) images. Here we follow the earlier classification
and introduce several new types of features documented by Venus Express.
Figure 13 shows examples of the Venus global views captured by the Venus Monitor-
ing Camera (VMC)/Venus Express (Titov et al. 2012) and UVI/Akatsuki camera (Yamazaki
et al. 2018). The VMC filter was centred at 0.365 μm, in the spectral band of the unknown
UV absorber. The Akatsuki camera took images at both 0.365 μmand0.283 μm, the latter
centred on the absorption features of SO2. Despite the large-scale similarity of morphology
patterns in two filters, there are certain differences in details (Limaye et al. 2018b). The radi-
ance at 0.283 μm is about ten times smaller than that measured at 0.365 μm. The 0.365 μm
image shows more contrast and a bright area in the equatorial region near the image centre.
This trend is the opposite to what was observed by OCPP/Pioneer Venus. Small-scale de-
tails in 0.283 μm images are muted or absent. Quantitative analysis of colocated 0.283 μm
and 0.365 μm (near) simultaneous images shows the correlation of brightness values at the
two wavelengths to be variable and part of the variation appears to be latitude dependent,
similar to what was found with Galileo and MESSENGER data at other wavelengths. The
UVI/Akatsuki show many features previously observed with contrasts decreasing near the
terminators.
Preliminary analysis of the IR (0.9μm) images of the day side revealed low-contrast
features whose appearance is quite different from that seen in UV (Limaye et al. 2018b)that
indicate variations of opacity on the deep cloud. The IR2 camera images taken at 2.02 μm,
the wavelength at which CO2absorption becomes significant, indicated brightness decrease
towards the poles that can be explained by deeper cloud top (Fig. 7, Ignatiev et al. 2009).
Figure 13 (lower panel) and Fig. 14 show equatorial views of the planet that clearly re-
veals the global “V” dark feature frequently observed by ground-based telescopes, Pioneer
Venus (Rossow et al. 1980) as well as during Galileo flyby. The figure emphasizes the re-
lation of the cloud pattern to latitude and local solar time. The morning sector is usually
covered with bright haze even in low latitudes. At the equator it takes zonal wind one Earth
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126 Page 18 of 61 D.V. Titov et al.
Fig. 13 Global views of Venus. Upper and middle panels: VMC UV (365 nm) images taken from a distance
of about 30,000 km with the sub-spacecraft point approximately in the middle latitudes of the Southern
hemisphere (Titov et al. 2012). The spatial resolution is about 25 km/px. The Southern pole is on the bottom
of the images. The atmosphere superrotates in the counter-clockwise direction (from right to left). Orbit
numbers are given at the bottom left of each image. The Venus Express orbital period was one Earth day.
Lower panel: examples of UVI/Akatsuki images in 0.283 μm(left)and0.365 μm (right) filters (Yamazaki
et al. 2018). North is at the top of the images
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Clouds and Hazes of Venus Page 19 of 61 126
Fig. 14 Dark “V” feature in the Venus UV images in a simple cylindrical projection: VMC/Venus Express
(left) (Titov et al. 2012) and SSI/Galileo (right) (Belton et al. 1991). The “sun” symbol in the left image
marks the sub-solar point. The atmosphere super-rotates from right (morning) to left (evening). Thick white
line shows the equator. Contours of the Earth’s continents are overplotted on the left image to illustrate
location and scale of the Venus global cloud features
day to bring air parcels to the sub-solar point. In the late morning the cloud tops become
darker suggesting that either dark material is brought from the depths of the cloud layer or
that the upper haze is destroyed (evaporated) by solar heating. Global streaks that extend
from the equator in the south-east direction usually form a sharp boundary between bright
high-latitudes and dark tropics. This boundary frames the afternoon convective region filled
with patchy clouds. The global cloud pattern suggests that the sub-solar point is some sort
of an “obstacle” for the flow deviating from purely zonal motion, the analogy also noticed
by Belton et al. (1976).
Venus low and middle latitudes (<50◦S) are generally darker than the high latitudes
(Figs. 13,14). The dark equatorial band is often present in the Venus tropics. Sometimes
the brightness minimum is displaced to the middle latitudes forming dark mid-latitude bands
separating slightly brighter equatorial region (circum-equatorial belt) and very bright high
latitudes. Mottled and patchy cloud patterns ubiquitously present at low latitudes suggest
significant role of convection and turbulence here (Titov et al. 2008), as prominently ob-
served by Mariner-10 (Belton et al. 1976), Pioneer Venus (Rossow et al. 1980) and Galileo
(Hueso and Sánchez-Lavega 2007). In the middle latitudes, the mottled clouds give way to
streaky features (bright streamers), indicating a transition to quasi-laminar flow at ∼50◦S.
The high latitudes are dominated by bright almost featureless cloud (bright polar band),
suggesting presence of a large amount of conservatively scattering aerosol that masks the
UV absorbers hidden in the cloud depth. Sometimes the boundary between the dark and the
bright regions is displaced to as low as 30◦S. The regions poleward from 70◦S are generally
slightly darker than the middle latitudes. They form a “polar cap” earlier observed by Pi-
oneer Venus at slant angles. Very often a narrow (few hundreds of kilometres) dark circle
appears at about 70◦S(Figs.13,17). All aspects of the planet’s appearance, including its
brightness, contrasts, and morphological pattern show strong variability on the time scale of
few days.
Figure 15 provides examples of a closer look at Venus low latitudes and shows the cloud
morphology in the Southern “tropics”. They cover low and middle latitudes from about
equator to the edge of the mid-latitude bright band (∼50◦S). The bright “lace” veil on top
of a darker cloud that extends to ∼30◦S is frequently present here. The images also show
bow shape waves in much more detail than documented by Pioneer-Venus (Rossow et al.
1980). The image taken in orbit 722 (Fig. 15) shows a pronounced afternoon convective
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126 Page 20 of 61 D.V. Titov et al.
Fig. 15 UV images of the Venus “tropics” captured by VMC/VEX from a distance of 10000–15000 km with
spatial resolution of 10–15 km/px (Titov et al. 2012). Lines and numbers mark latitudes. Image centres are
close to the local noon. Orbit numbers are given in the bottom left
wake featuring well-developed turbulence downstream of the sub-solar point. The image
from orbit 920 shows the sub-solar point dominated by dark clouds with small wind streaks
suggesting flow diverging from the sub-solar point.
Mesoscale images in Fig. 16 zoom in on the transition to the bright mid-latitude band.
The remarkable change usually occurs at 50◦–60◦S(Fig.13) and occupies about 1000 km
(∼10◦latitude). The region is characterised by drastic change from mottled to streaky cloud
morphology. This trend in the global cloud pattern implies that convection vanishes pole-
ward of ∼50◦S that is consistent with the convectively stable temperature structure in the
“cold collar” region (Limaye et al. 2018a). Thin cloud streaks, thousands of kilometres long,
are typical for this zone, implying that quasi-laminar flow completely dominates over tur-
bulent mixing. The streaks are tilted with respect to the latitude circles (see also Fig. 14)
indicating a relation between zonal and meridional wind components. As pointed out by
Schinder et al. (1990) they could be formed by a dominant zonal motion combined with a
poleward advection and shearing of the clouds by the winds with possible action of a super-
imposed wave. The ratio of the mean meridional and zonal wind components can be assessed
from the streaks’ slope (Fig. 14) and is in agreement with zonal and meridional wind veloci-
ties of approximately 100 m/s and 15 m/s derived from the cloud tracking (Sánchez-Lavega
et al. 2017).
The long streaks with sharp brightness contrasts in the transition region suggest a strong
jet stream flowing along the equatorward edge of the bright mid-latitude band. The position
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Clouds and Hazes of Venus Page 21 of 61 126
Fig. 16 UV images of the mid-latitude transition region taken by VMC/VEX from a distance of
10000–15000 km with resolution of 10–15 km/px (Titov et al. 2012). The South pole is on the bottom.
Orbit numbers are given in the lower left of each image
Fig. 17 UV images of the Southern polar region from a distance of about 15000 km and spatial resolution
of ∼10 km/px (Titov et al. 2012). The pole is in the bottom. Orbit numbers are given in the bottom left of
each image
of this morphological feature is roughly consistent with the mid-latitude jet observed in the
zonal wind field derived from the cloud tracking (Khatuntsev et al. 2013) and calculated
from the temperature field in cyclostrophic approximation (Piccialli et al. 2011). The jet
roughly follows the equatorial edge of the “cold collar” region (Limaye et al. 2018a).
Figure 17 shows examples of polar images. Poleward of 50–60◦S the clouds become
very bright and uniform, suggesting that the UV absorbers are either absent here or more
likely hidden deep below the cloud top overlaid by a thick sulphuric acid hood. This may be
due to suppression of the convection that brings absorbers from the deep cloud (Titov et al.
2008). In the periods of reduced polar hood the polar region is often crossed by thin dark
circular or spiral “grooves” that are a few hundred kilometres wide and are likely created
by local jets. The most frequently observed feature is the dark polar oval located at ∼70◦S
(Fig. 17). This almost axisymmetric structure sometimes disappears and dark features in the
“polar cap” are distributed chaotically (see Titov et al. 2012). Appearance of the “polar cap”
strongly varies. For instance, in orbit 1015 (Fig. 17) the polar hood was very thick, uniform
and extended to 40◦S, while in orbit 1259 the darker main cloud deck is clearly visible.
Close-up images of the Northern hemisphere reveal many features indicating convective
activity, turbulence and waves at the cloud tops (Fig. 18). At a spatial resolution of few kilo-
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126 Page 22 of 61 D.V. Titov et al.
Fig. 18 Small scale features at
the Venus cloud top: UV images
of the low latitudes (upper panel
and lower left image) and near IR
image (lower right) taken from a
distance of 3000–5000 km with
spatial resolution of few
kilometres per pixel (Titov et al.
2012). White bars in the lower
right of each image show the
scale distance of 200 km. Orbit
numbers are given at the bottom
left of each image
metres the cloud top has patchy morphology with dark spots and “valleys” a few tens of
kilometres in size (orbits 269, 590). The dark material appears to be hidden in the depths
of the cloud and becomes visible through openings in the upper cloud. The mottled clouds
form convective cells with typical size of 100–200 km and “wave trains” with similar wave-
length (orbit 269). In some cases they resemble Earth cumulus cloud columns a few tens of
kilometres across (orbits 150, 590). These clouds show signs of lateral advection indicating
strong wind shear at the visible cloud tops (orbit 150) as originally suggested by Crisp and
Young (1978).
At high latitudes (>60◦N) three types of waves: long straight features, short wave trains,
and irregular wave fields—are often observed. The long waves (bottom right in Fig. 18)have
wavelengths of a few tens of kilometres and extend for a few hundred kilometres. Short
waves form compact “trains” several tens of kilometres wide with typical wavelengths of
3–7 km. The trains often originate at the fronts of long features and seem to be genetically
related to them. Irregular wave fields consist of chaotically distributed features with a size of
several kilometres. Such irregular wave fields often overlap with short regular waves. Inter-
estingly, the waves are seen in all VMC channels, suggesting that their origin is not related
to inhomogeneities in the near-UV absorbers distribution, but they are rather produced by
variations in haze opacity or, more likely, changes in the solar illumination angle across the
wave (Piccialli et al. 2014; Sánchez-Lavega et al. 2017). We note that the cloud tops in the
VMC images at 0.965 μm is located at 61–67 km close to that in UV and near-IR (1.5μm)
(see Sect. 2and Fig. 6).
The observed global cloud pattern strongly supports the idea of a planet scale vortex
circulation first discovered by UV imaging (Suomi and Limaye 1978).
VIRTIS/Venus Express observations in the thermal IR range (3–5 μm) (Piccioni et al.
2007) significantly contributed to our understanding of the polar cloud morphology and
dynamics by providing details that complement the Pioneer Venus observations (Taylor et al.
1980). The most remarkable feature observed by Venus Express “in action” is the polar eye
of the planetary vortex. Figure 19 shows examples of its view at 5 μm.
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Clouds and Hazes of Venus Page 23 of 61 126
Fig. 19 Morphology of the “eye” of the planetary vortex: left—sketch of the vortex “eye” (red) overplotted
on top of a VMC UV image (grey), middle and right—original images of the polar “eye” taken at 5 μm. The
colour approximately represents the cloud top temperature
The brightness pattern at thermal IR wavelengths reveals spatial variations of the cloud
top temperature. The polar “eye” is usually confined within ∼70◦S latitude circle thus hav-
ing the size of few thousands of kilometres. It is located within the dark polar oval in the
UV images (Figs. 12,17; Titov et al. 2008). The spiral arms of the IR features are often
connected to the dark oval. In some cases the vortex eye or its spiral arms have counterparts
in the simultaneously captured UV images. All of this suggests that the UV and thermal
IR features observed at the pole are both manifestation of the same dynamical processes
in the polar atmosphere and that the global UV pattern is likely to be created by the tem-
perature and dynamical conditions at the cloud tops (Titov et al. 2008). The polar “eye”
has a very variable appearance spanning from a simple oval to complex multi-pole struc-
tures (Fig. 19). It has remarkable morphological similarities to Earth’s tropical cyclones and
hurricanes, with vivid dynamics. Limaye et al. (2009) succeeded to simulate the “eye” mor-
phology with an idealized nonlinear and non-divergent barotropic model and found similar
structures in the modelled vorticity field. Luz et al. (2011) showed that the centre of rotation
of the polar “eye” is displaced from the South pole by typically ∼3 degrees of latitude and
drifts around the pole.
One of the first major results from the LIR/Akatsuki camera was discovery of a
bright band aligned almost north–south with a slight curvature which has been inter-
preted as a standing gravity wave triggered by surface topography (Fukuhara et al. 2017;
Navarro et al. 2018; Limaye et al. 2018b). The feature had brightness temperature contrast
of ∼5K and lasted as long as 2–3 weeks. Surprisingly, the signature of the standing wave
can also be detected in UVI and IR2 (2.02 μm) dayside images (Satoh et al. 2017). The
detection at 2.02 μm suggests that cloud-top variations likely occur due to the wave. How-
ever the stationary wave has not yet been detected either in 0.9-μm daytime IR1/Akatsuki
images, or MDIS/MESSENGER images or VIRTIS/Venus Express data that could probably
indicate transient nature of the wave. Observations in the near-IR transparency “windows”
on the night side enabled sounding of the deep cloud morphology that is not visible from
orbit at other wavelengths. In this case the cloud layer is illuminated from below by ther-
mal emission from the hot surface and the lower atmosphere. The features revealed by these
images are produced by the cloud opacity variations occurring mainly in the middle and
lower cloud layers (50–60 km). Figure 12 (lower left) shows equatorial view of Venus in
the 2.3μm “window” captured by NIMS/Galileo (Carlson et al. 1993). This image reveals
patchy morphology of the deep cloud at low latitudes. Recently IR2 camera onboard the
Akatsuki spacecraft (Satoh et al. 2017) captured views of the planet night side in the near-
IR spectral “windows” revealing morphology of the deep cloud (Fig. 20) (Limaye et al.
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126 Page 24 of 61 D.V. Titov et al.
Fig. 20 Observations of the deep cloud morphology in the 2.3μm “window on the night side: left—averaged
radiation field from VIRTIS/VEX (Cardesín-Moinelo et al. 2010). South pole is in the middle; right-false
colour equatorial view by IR2/Akatsuki camera (credit JAXA/ISAS/DARTS/Damia Bouic)
2018b). The low latitudes are full of small- and large-scale features: waves, mushroom-
shaped mesoscale vortices, long linear streaks, sharp-contrast fronts etc. Some images show
high-contrast sharp linear boundaries running roughly north–south that can be explained by
strong opacity variations in adjacent air masses echoing the results Vega balloon (Lorenz
et al. 2018). However, the origin of different air masses is still an open question. The cloud
morphology in high latitudes is smoother and streaky resembling that of the cloud tops seen
in UV (Figs. 13–17). The nightside images suggest that dynamics plays significant role in
their formation. Moreover, they might indicate a mixture of unknown dynamics as well as
compositional differences that must be responsible for the differences in the appearance of
the features.
The polar view of the planet (Fig. 20) at the same wavelength indicates that the emis-
sion reaching space from the lower atmosphere reaches its maximum in the middle lati-
tudes (40–60◦S). The emission slightly decreases towards the equator that could be due to
higher observation angle. The emission abruptly drops poleward of ∼60◦latitude (Cardesín-
Moinelo et al. 2010). This feature can be explained either by supposing that the cloud opacity
at 1 μm increases by a factor of 2–3 reaching values of 60–90 in the polar regions or that the
cloud composition in the polar region differs from sulphuric acid. There are also evidences of
anomalously large particles present in the deep cloud in the polar regions (Wilson et al. 2008;
Barstow et al. 2011).
The latitude of the emission abrupt drop coincides with the “cold collar” region with the
coldest temperature at the cloud top (Limaye et al. 2018a) and is located at the poleward
side of the midlatitude jet (Sánchez-Lavega et al. 2017). Figure 20 suggests that the cloud
properties in the polar and middle latitudes are quite different with a sharp separation at
60◦S and suppressed material exchange between the regions. Analogous transport “barriers”
were observed on other planets resulting in strong composition gradients in the stratospheres
of Titan (Teanby et al. 2008) and Earth (polar ozone “hole”). Such mixing “barriers” are
associated with regions of maximum gradient of potential vorticity that seems to be true
also for Venus (Sánchez-Lavega et al. 2017).
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Clouds and Hazes of Venus Page 25 of 61 126
Fig. 21 Correlation of the global morphology patterns in the upper cloud (grey UV images) and in the deep
cloud (red-orange images at 1.7μm on the night side)
Simultaneous Venus Express observations in UV and near-IR transparency “window” on
the night side (Fig. 21) show remarkable correlation of the global cloud morphology at the
cloud tops and in the middle cloud.
4 Microphysical Properties of the Cloud Population
The upper haze was known to be composed of submicron aerosol particles with an effec-
tive radius of about 0.25–0.29 μm, also present within the clouds. The refractive index of
1.435–1.45 at 0.55 μm is in agreement with sulphuric acid composition of the haze particles
(Hansen and Hovenier 1974; Kawabata et al. 1980; Krasnopolsky 1983;Satoetal.1996;
Braak et al. 2002). The progress in this field achieved after publication of the Venus and
Venus-II books (Esposito et al. 1983,1997) is mainly due to remote sensing observations.
Properties of the upper haze and upper cloud were revealed by solar occultation. Deep at-
mospheric sounding in the near-IR transparency “windows” provided important insights in
the microphysics of the middle and lower cloud.
Microphysical properties of the mesospheric haze were inferred from the solar occulta-
tion in UV through the near-IR spectral range performed by SPICAV-SOIR onboard Venus
Express. Wilquet et al. (2009) derived the wavelength dependence of aerosol extinction re-
lated to the effective radius and composition of the particles (Fig. 22). Comparison of the
spectral dependence of extinction coefficient derived from the measurements to that calcu-
lated for mode 1 aerosol distribution (Table 1) suggests that either the assumption on the
upper haze sulphuric acid composition or size distribution need to be revisited. The study
demonstrated for the first time existence in some cases of at least two types of particles:
submicron haze with r=0.1–0.3μm and larger particles with r=0.4–1 μm. Therefore, the
model describing the upper haze on Venus should include a bimodal population (Wilquet
et al. 2009).
Luginin et al. (2016) analysed about 200 solar occultations obtained by SPICAV-IR on-
board Venus Express. A bimodal distribution was found to be typical for 75–85 km altitude
range while unimodal particle population dominates at 70–75 km and above 85 km. Table 2
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126 Page 26 of 61 D.V. Titov et al.
Fig. 22 Spectral dependence of
aerosol extinction derived from
SPICAV-SOIR spectra (brown
line and circles) and normalized
extinction calculated from the
Mie theory assuming 75%
H2SO4aerosols and a unimodal
log-normal size distribution with
reff =0.3μmandνeff =0.18
(cyan solid line)
Tab le 2 Latitudinal variation of the mean effective radius. Reff corresponds to the unimodal distribution;
reff1 and reff2 correspond to modes 1 and 2 of the bimodal case
Altitude, km 60◦S–60◦N60
◦N–90◦N
Reff,μmreff1,μmreff2 ,μmReff,μmreff1,μmreff2,μm
>85 0.42 ±0.14 0.12 ±0.02 0.78 ±0.12 0.30 ±0.14 0.12 ±0.05 0.81 ±0.08
80–85 0.75 ±0.09 0.12 ±0.02 0.82 ±0.07 0.32 ±0.15 0.13 ±0.03 0.81 ±0.11
75–80 0.79 ±0.06 0.11 ±0.03 0.89 ±0.08 0.47 ±0.20 0.12 ±0.04 0.82 ±0.12
70–75 0.77 ±0.07 0.13 ±0.03 0.97 ±0.17 0.59 ±0.18 0.11 ±0.07 0.82 ±0.12
summarises the results. For the unimodal size distribution, the effective radius is larger in
the middle and equatorial latitudes than that observed near the North Pole. No statistically
significant differences were found in the size distribution between morning and evening.
Figure 23 shows variability of the retrieved size distribution over the mission. When a uni-
modal size distribution is sufficient to reproduce the aerosol extinction, the effective radius
(reff) varies greatly between 0.2μmand1.0μm, while for the bimodal distribution the
ranges of values for reff1 and reff2 are much smaller.
Luginin et al. (2016) derived particle number density profiles for modes 1 and 2 as shown
in Fig. 24. The population of both modes gradually decreases with altitude from ∼500 cm−3
at 75 km to ∼50 cm−3at 90 km for mode 1, and from ∼1cm
−3at 75 km to ∼0.1cm
−3at
90 km for mode 2. The ratio between mode 1 and 2 number densities is ∼500 and does not
change with altitude. The mean parameters of mode 1 agree with the haze properties summa-
rized in the first edition of the Venus book (Esposito et al. 1983) wherein the number density
at 70–90 km is ∼500 cm−3. The most significant difference with earlier investigations is
detection of larger particles (r>0.2μm) ubiquitously present at 80–90 km.
Limb imaging in thermal IR range by VIRTIS/VEX confirmed the presence of micron
size sulphuric acid particles in the upper haze. de Kok et al. (2011) analysed nightside spec-
tra at 4.5–5.0μm of the Venus limb. Assuming log-normal size distribution and 75% H2SO4
particle composition the authors retrieved vertical profiles of 1 μm-sized mode 2 particles.
The number density monotonically decreased with height demonstrating higher variability
at middle latitudes than at low latitudes. These results are in general agreement with those
derived from the solar occultation by SPICAV-IR and SOIR at high northern latitudes (Wil-
quet et al. 2009). These works are therefore complementary in terms of local solar time
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Clouds and Hazes of Venus Page 27 of 61 126
Fig. 23 The retrieved mean effective radius for four altitude ranges over the mission. Each occultation mea-
surement is averaged within the altitude bin. Green and red symbols correspond to mode 1 and mode 2 of the
bimodal distribution. The blue symbols represent unimodal distribution. Big crosses show mean values. The
x-axis is the orbit number for the upper panels and the corresponding date for the bottom panels of the figure
(Luginin et al. 2016)
(nightside vs terminator) and latitudes (equatorial and middle vs. polar latitudes). They also
suggest that the transition between the upper cloud and overlaying haze is less abrupt than
was previously thought.
Imaging at varying phase angle enabled sounding of both mesospheric haze and the cloud
top properties. Markiewicz et al. (2014) discovered a glory in the VMC/VEX observations of
the planet at small phase angles in three wavelengths of 0.365 μm, 0.513 μmand0.965 μm.
Glory is an optical phenomenon that poses stringent constraints on the cloud properties. The
very fact that the glory was observed implies that the scattering medium is rather homo-
geneous and consists of spherical particles with narrow size distribution. From the angular
position of the glory features Petrova et al. (2015) and Markiewicz et al. (2018) estimated
particle effective radius reff =1.0–1.4μm for different regions of the cloud deck that cor-
responds to the mode 2 of the particle size distribution. Interestingly, some features of the
glory implied a real part of the refractive index higher than that of concentrated sulphuric
acid. The authors suggested this material to be ferric chloride or sulphur, both of which are
candidates for the unknown UV-blue absorber in the upper cloud layer of Venus. The change
of the glory maximum position in the UV phase curve indicated decrease of the particle size
from 1.05 μmto0.8–0.9μm with latitude (40◦S–60◦S) before local noon. Shalygina et al.
(2015) analysed the full set of VMC observation at 965 nm. The results showed tempo-
ral and spatial variations of the cloud properties. In general, the particles at low latitudes
were larger than those in the southern polar regions (reff =1.2–1.4μmvs. 0.9–1.05 μm). At
40◦S–60◦S the refractive index was usually smaller than that in the other regions (1.44–1.45
vs. 1.45–1.47). Small submicron (reff ∼0.23 μm) particles are detected mostly in the morn-
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126 Page 28 of 61 D.V. Titov et al.
Fig. 24 Mean profiles of the
particle number density based on
the analysis of SPICAV-IR/VEX
solar occultation (Luginin et al.
2016) for mode 1 (green) and
mode 2 (red)
ing. Rossi et al. (2015) analysed degree of polarization at low phase angles measured by
SPICAV in the 0.65–1.7μm range. The results are consistent with mean values of effective
radius reff ∼1μm and its variance νeff ∼0.07 and a refractive index nr=1.42 ±0.02 at
λ=1.1μm in good agreement with previous determinations.
In 2016 JAXA’s Akatsuki spacecraft had started regular orbital observations. Analysis of
the phase curves acquired by IR1 (∼1μm) and IR2 (∼2μm) cameras on board Akatsuki
during the first attempt of orbit insertion in 2010 already suggested existence of 1 μm parti-
cles at the cloud top. Satoh et al. (2015) had to modify the standard cloud model (Esposito
et al. 1983) that fits both Pioneer Venus and recent ground based observations (Mallama et al.
2006; García Muñoz et al. 2014) extending population of micron-size particles to higher al-
titudes (up to 75 km) and adding large mode 3 particles in the upper cloud to reproduce
the observed phase curve at low phase angles. Lee et al. (2017) derived reff =1.26 μmand
variance νeff =0.076 from the phase angle dependence of the planet global albedo derived
from the images captured by Akatsuki UVI camera at 0.283 μmand0.365 μm in agreement
with the earlier results on mode 2 particles in the upper cloud.
Spectroscopy in the near-IR transparency “windows” on the night side provided new
information about global distribution of the microphysical properties of the deep cloud. The
method exploits correlation between the radiances measured at 1.74 μmand2.3μmwhich
are sensitive to the aerosol properties, in particular, proportion between mode 2’ and mode 3
(Table 1) and concentration of sulphuric acid in the deep cloud particles. Analysing the
NIMS/Galileo data Carlson et al. (1993) favoured the conclusion about significant spatial
variability of particle population in the deep cloud with a tendency for larger particles in
the Northern hemisphere. The analysis by Grinspoon et al. (1993) indicated that opacity at
1μm ranged from 25 to 40, reinforcing the conclusion that the “typical” cloud properties
measured by Pioneer Venus (Table 1) should be adopted with caution.
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Clouds and Hazes of Venus Page 29 of 61 126
A similar analysis of near-infrared emissions on the nightside of Venus observed by
VIRTIS/Venus Express revealed anomalous cloud particles in the polar regions (Wilson et al.
2008). These particles were found close to the centers of the polar vortices at both poles and
are either larger or differ in composition from those elsewhere in the planet. This result was
confirmed by Barstow et al. (2011) who also found that the acid concentration increases with
opacity resulting in 90–100% concentration in the polar regions. A comprehensive analysis
by Haus et al. (2013) showed that the cloud particle size and total opacity exhibit a minimum
at 50◦N and increase towards the equator and the North pole.
5 Cloud Composition
The composition and chemistry of Venus clouds was reviewed in both Venus and Venus-2
books (Esposito et al. 1983,1997). The Venus clouds are composed primarily of liquid
sulphuric acid mixed with water. However, there must be other constituents as well. It is
also not clear whether cloud condensation nuclei play an important role on Venus, and if
so what their composition is. Particulates have been detected at altitudes below the main
cloud base at 48 km, where temperatures are too high to allow sulphuric acid droplets to
exist. And finally, the unknown UV-blue absorbers responsible for absorption at 0.3–0.5μm
in the upper cloud are likely to be particulate species. In this section we will first review
the evidence for sulphuric acid and then we will discuss evidence for particulates of other
composition. Gaseous chemical cycles relevant to the formation of aerosols are reviewed in
a dedicated chapter of this Book (Marcq et al. 2017).
Sulphuric Acid The first indications that the clouds of Venus were composed of sulphuric
acid droplets came from ground-based observations. Polarimetric observations of Venus
were shown to be consistent with scattering from spherical droplets with a narrow size dis-
tribution with a mean radius ∼1μm and a refractive index of 1.44 ±0.02 at λ=0.55 μm
(Hansen and Hovenier 1974) that was attributed to ∼75% by weight sulphuric acid. Spec-
tral absorption features near 3 μm, 11.2μm and 25 μm seen in Venus thermal infrared emis-
sion spectra were found to match the spectrum of sulphuric acid (Young and Young 1973;
Zasova et al. 2007). Finally, the altitude of the cloud-base, as revealed by descent probes,
was found to lie at around 48 km, at a temperature consistent with that expected for thermal
decomposition of sulphuric acid.
In situ measurements of particle composition confirmed that the clouds are composed
primarily of sulphuric acid mixed with water. The Pioneer Venus mass spectrometer did not
have a dedicated aerosol sampling inlet, but it did experience blockages attributed to aerosol
droplets getting stuck in its inlet. When the droplets evaporated, the evolved gases were
consistent with those which would be expected from a droplet composed of 85% H2SO4and
15% H2O(Hoffmanetal.1980). The most direct evidence of sulphuric acid composition of
cloud droplets was finally provided by the Vega probes. Atmosphere was pumped through
aerosol collecting filters, so that the composition of the aerosol droplets could be determined
separately from that of the gaseous species. Cloud droplets were collected on carbon fibres
while the probe descended from 63 km to 48 km, and then the sample was sealed off and
heated. The evolved gases were analysed by gas chromatography and mass spectrometry and
were found to consist of SO2,H
2O, and CO2as expected from sulphuric acid droplets heated
on a carbon fibre substrate (Gel’man et al. 1986). In addition to sulphuric acid the Vega-1
X-ray fluorescent spectrometer revealed significant amounts of chlorine and phosphorus in
the cloud droplets, as will be discussed below.
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126 Page 30 of 61 D.V. Titov et al.
It is from scattered light sensors (particle size spectrometers and nephelometers) on Pio-
neer Venus, Venera and Vega probes that the cloud vertical structure and particle size distri-
butions were determined (Fig. 1). Nephelometry does not provide a strong constraint on the
refractive index (and thus on the composition) of small particles (mode 0/mode 1), because
their phase function is close to that of Rayleigh scatterers. The modes 2 and 2’ particles were
found to have refractive indices of 1.35–1.50 as summarised in Knollenberg and Hunten
(1980) and Ragent et al. (1985). This range includes the value of 1.42–1.47 expected for
sulphuric acid solution at visible wavelengths at temperatures of 200–350 K. Initial analy-
ses of LCPS data showed a population of large non-spherical mode 3 particles (Knollenberg
and Hunten 1979), but subsequent studies suggested that the mode 3 could in fact be simply
a large “tail” of the liquid mode 2’ distribution, once calibration errors are taken into account
(Toon et al. 1984). The presence of mode 3 and in general particles with r>5μmalsowas
not confirmed by the particle size spectrometers onboard Vega descent probes (Zasova et al.
1996).
There have not been further in situ measurements in the Venus clouds since the Vega
probes in 1985. However, a large number of remote sensing observations have been
provided by Venus Express, Galileo and ground-based telescopes. The most spectacular
new observations, perhaps, are those of ‘glories’—rainbow-like circular patterns imaged
from Venus Express’s VMC camera on the cloud tops. Analysis of these features finds
cloud top refractive index to vary widely over the range 1.44–1.53 (Petrova et al. 2015;
Markiewicz et al. 2018). The higher end of this refractive index range is too high to be ex-
plained by sulphuric acid/water mixtures. No consistent correlations of the refractive index
with UV-blue albedo have been reported to date, but analyses are ongoing. Initial analyses
of polarimetric observations by SPICAV suggests lower refractive indices of 1.42 ±0.01 at
λ=1.101 μm (Rossi et al. 2015) but these results are somewhat preliminary and may be
sensitive to aerosol at altitudes different from those probed by the glories.
Further indirect constraints on particle composition in the lower and middle clouds can
also be obtained from observations in the near-infrared spectral “windows” at 1.7μmand
2.3μm on the nightside. As has been discussed, the emission in these spectral domains
is sensitive to cloud thickness, and the relative brightness in the different spectral bands
within the transparency “windows” depends on cloud particle size and refractive index. If
the clouds are modelled solely as combinations of modes 1, 2, 2’ and 3 of sulphuric acid
droplet populations, then the band ratios can be best matched by having particles of 80%
H2SO4: 20% H2O in optically thin clouds, rising to 95% H2SO4:5%H
2O in regions of
thick cloud (Barstow et al. 2011). The same analysis also found a slight increase of H2SO4
concentration from ∼84% H2SO4at low latitudes to ∼90% at high latitudes. It must be
emphasized, though, that there are many other parameters to which these band ratios are
sensitive, in particular the size distributions of the cloud particles. A marked change in the
brightness ratio of the 1.7/2.3μm “windows” at the core of the polar vortex is seen, which
may be associated with changing particle sizes, but could equally signal a change in particle
composition (Wilson et al. 2008).
Recent ground-based observations supported space missions in the study of Venus clouds
and their relation to the minor species. Arney et al. (2014) mapped emissions in the near-
IR spectral “windows” on the Venus night side using the Apache Point Observatory 3.5 m
telescope TripleSpec spectrograph. They produced the first simultaneous maps of water va-
por, CO, HCl, OCS, SO2below the clouds and sulphuric acid concentration and total cloud
opacity and established some trends between the cloud properties and gaseous species. The
retrieved mean cloud opacity at ∼2.0μm was found to be about 20, i.e. by a factor of 1.7–
1.8 higher than that in the post-Pioneer Venus model by Crisp (1986). The cloud opacity
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Clouds and Hazes of Venus Page 31 of 61 126
is highly variable but the features with the most opaque clouds tend to be near the equator.
The average value of sulphuric acid concentration was 79±4%—slightly higher than in the
earlier models (75%)—with a trend to higher concentrations in the Northern hemisphere.
The authors also tentatively confirmed the conclusion by Barstow et al. (2011) that the acid
percentage correlates with the total opacity.
Analysis of the VIRTIS/Venus Express observations (Tsang et al. 2010; Barstow et al.
2011) suggested an anticorrelation between the cloud opacity and water vapor mixing ra-
tio at 30–45 km. This may be due to breakdown of cloud particles below the cloud deck
(Krasnopolsky 2007). In this case, optically thin clouds indicate cloud dissolution and there-
fore enhanced water. The correlations between the cloud opacity and abundance of minor
species observed by Arney et al. (2014) were less pronounced. The authors reported evi-
dence of correlations between the cloud deck and several gases, suggesting active processes
in the Venus lower atmosphere. They also observed a marked hemispherical dichotomy in
sulfuric acid, water vapor, and sulfur dioxide. All three parameters were enhanced in the
Northern hemisphere. Sulfuric acid virga, the rain that evaporates before reaching the ground
(Gao et al. 2014), was suggested as a possible mechanism for the water vapor dichotomy.
Alternatively the dichotomy can be related to peculiarities of the atmospheric dynamics.
We note however that ground-based observations in the 2.3μm “window” by Marcq et al.
(2006) did not indicate significant latitudinal variability of the subcloud water abundance.
Unknown UV-Blue Absorbers The unknown UV-blue absorbers responsible for the
striking contrasts seen in the images such as Figs. 13–18 cannot be explained by clouds
composed solely of sulphuric acid. Some absorption below 0.3μm is due to gaseous SO2.
Spectral behaviour of the planet albedo particularly in the 0.3–0.4μm region requires pres-
ence of an additional absorber. The absence of sharp spectral features in its spectrum, as
well as results from Venera and Pioneer Venus photometers, indicate that the unknown
absorption is due to a particulate rather than a gaseous species (Ekonomov et al. 1984;
Tomasko et al. 1980) and occurs mainly above 57 km. A vast number of candidates for this
UV absorber have been proposed, from sulphur and chlorine species to organic compounds
and mineral dusts. Recent reviews of atmospheric composition and chemistry as well as
candidate materials can be found in Mills et al. (2007), Zhang et al. (2012)andMarcqetal.
(2017).
Elemental sulphur was proposed as a candidate for the unknown absorber by Hapke
and Nelson (1975). However elemental sulphur and its S3and S4allotropes fail to fit the
observed near-UV absorption. Recently Carlson et al. (2016) found that some of sulphur
allotropes could form aerosols at cloud level on Venus producing a good spectral match to
the Venus albedo. However, concerns have been raised over whether the vertical distribution
of polyatomic sulphur would match that of the unknown UV absorber. All photochemical
models since Yung and DeMore (1982) predicted negligible abundances of aerosol sulphur
in the upper cloud layer. According to the chemical kinetic model (Krasnopolsky 2013),
significant abundances of gaseous sulphur are produced in the lower atmosphere below the
clouds. This sulphur should condense near 48 km and form an aerosol layer (Krasnopolsky
2016) with mass loading of ≈10% of that in the lower cloud layer observed by the Pioneer
Venus particle size spectrometer. This proportion is similar to the sulphur-to-sulphuric acid
aerosol mass loading ratios of 1:10 observed by the Vega nephelometer and particle size
spectrometer (Gnedykh et al. 1987) and of 1:7 observed by the Vega gas chromatograph
(Porshnev et al. 1987). Furthermore, microphysical models of the Venus clouds suggest
that soluble condensation nuclei are required for realistic cloud formation, but polyatomic
sulphur is not soluble in sulphuric acid and so other species would be needed.
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126 Page 32 of 61 D.V. Titov et al.
There have been several attempts to constrain the composition of the UV-blue absorbers
by finding correlations between observed cloud morphology in UV images and atmospheric
composition, but to date these have been largely inconclusive. Water vapour above the clouds
was mapped by Venus Express using both the SPICAV and VIRTIS spectrometers (Fedorova
et al. 2016; Cottini et al. 2012,2015). No consistent correlations were found for small-scale
variations between UV features and abundances of water vapour or sulphur dioxide. On a
larger scale, both works show that water concentrations are elevated at latitudes of ±60–80◦,
that is likely related to the several kilometres lower position of the cloud top.
Analysis of the photometrically corrected VMC/VEX UV images by Lee et al. (2015)
revealed a decline of the mean albedo at 0.365 μm by 25–40% over 2000 orbits in 2006
through 2011 that could be partly due to slow decrease of the UV channel sensitivity. The
relative brightness contrast between high and low latitudes was also observed to change and
this cannot be due to detector degradation. Lee et al. (2015) identified short-term contrast
variations of up to 40% correlated with mesospheric SO2abundance. Since the global con-
trast variability was mainly due to albedo changes at high latitudes, the authors explained
the correlation by intensive formation of sulphuric acid haze in the periods of mesospheric
SO2enhancement. Lee et al. (2015) also found that the UV contrast slowly decreased over
almost 6 years of Venus Express observations, apparently correlated with the long-term de-
crease of mesospheric SO2abundance. The observed increase of the contrasts phase angle
dependence possibly implies variations of the vertical motions at the cloud top level on
meso- or global scale.
The images recently taken by Akatsuki UVI camera at 0.283 μmand0.365 μmshowed
morphological differences suggesting that the spatial distributions of SO2and unknown UV-
blue absorbers sounded respectively at these wavelengths are governed by, at least partly,
different chemical and/or dynamical processes (Yamazaki et al. 2018).
Observations at UV-blue wavelengths by VMC and VIRTIS onboard Venus Express pro-
vided the best opportunity to study spatial and temporal distribution of the unknown ab-
sorbers since the Pioneer Venus mission. Molaverdikhani et al. (2012) found two models of
the UV absorber vertical distribution that match the VMC data set. One model is a well-
mixed absorber above ∼63 km that is consistent with the altitude of photochemical for-
mation of sulphuric acid. The second model suggests a thin layer of pure absorber placed
roughly around 71 km, the altitude consistent with the cloud top. This echoed previous work
by Braak et al. (2002) using Pioneer Venus OCPP data, which similarly found that a wide
range of vertical distribution models could satisfactorily fit the observations. Molaverdikhani
et al. (2012) concluded that the average optical depth of the absorber at 0.365 μminthe
equatorial region is 0.21 ±0.04 decreasing to τ=0.08 ±0.05 towards the pole.
Petrova et al. (2015) examined the phase function of dark and bright regions at both UV
and near-IR wavelengths at low phase angles in VMC/VEX images. The authors noted that
the contrasts occurred usually in the UV channel only, implying that the UV absorber is
present only in small (∼0.3μm) particles. The contrasts are sometimes observed also in the
near-IR wavelengths; this would be consistent with UV absorber being present also in 1 μm
particles which are the dominant population in the upper cloud. They suggested that the UV
absorber could be present in 0.3μm aerosol mode, which acts as a condensation nucleus for
larger cloud particles.
Titov et al. (2008) found correlation of UV-dark areas with the regions of enhanced con-
vective mixing in the upper cloud inferred from the global cloud morphology pattern and
dynamical properties. Similarly, Bertaux et al. (2016) concluded that regions near Aphrodite
Terra were on average UV-dark and exhibited high water vapour concentrations. Both ab-
sence of correlation between the UV absorber and mesospheric water vapour abundance
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Clouds and Hazes of Venus Page 33 of 61 126
and morphological evidences support the hypothesis that the UV-blue absorbers are brought
to the cloud tops from depth, along with tropospheric water vapour, with no evidence that
water vapour constrains formation of the UV absorber. Simultaneous VMC imaging in the
UV and visible filters tentatively indicated that the absorption is likely to extend to visible
wavelengths (Titov et al. 2012).
Most of the researchers associate the unknown UV-blue absorption with cloud particles.
However, there are recent attempts to suggest gaseous species as candidates. Frandsen et al.
(2016) studied chemistry and optical properties of S2O2and found that this species can
form in a sulphur cycle at the cloud tops and that its spectral properties match those of the
unknown near-UV absorber reasonably well. The problem of this identification is that the
photochemical lifetime of S2O2is very short (few seconds in sunlight) that precludes its
existence on the day side of the planet. Krasnopolsky (2018) included the S2O2formation
rate and absorption cross sections from Frandsen et al. (2016) in his photochemical model.
The predicted S2O2absorption was found to be about two orders of magnitude weaker than
that observed.
Iron was detected in cloud droplets by the Venera-14 X-ray fluorescent spectrometer
(Petryanov et al. 1981). Zasova et al. (1981) proposed ∼1% FeCl3solution in sulphuric
acid mode 2 particles as the near-UV absorber. Recently this idea was elaborated in detail
by Krasnopolsky (2017) summarising chemical, dynamical and spectroscopic arguments in
favour of this species being additional constituent in the Venus clouds responsible for the
UV-blue absorption. The analyses confirmed that FeCl3vertical distribution matches that
measured by descent probes, and that it is soluble in sulphuric acid so could act as a con-
densation nuclei. The resulting solution would not only have a refractive index higher than
that of sulphuric acid, as required to match the glory observations (Petrova et al. 2015;
Markiewicz et al. 2018), but also would absorb light at 0.32–0.5μm explaining the UV-blue
absorption in the Venus spectrum. If iron chloride is present in the rocks with abundance of
19 ppb, its vapor would condense at 47.5 km and form mode 1 aerosol in the lower and mid-
dle cloud layers. These particles can become condensation nuclei in the upper cloud layer,
and mode 1/mode 2 mass ratio and vertical flux would be in agreement with the required
concentration of 1%. The solution is rather stable at low temperatures and converts to col-
orless iron sulphate at room temperature (<58 km) that explains vanishing of absorption
detected by Venera-14 photometer (Ekonomov et al. 1984).
Pérez-Hoyos et al. (2018) analysed the spectra from 0.3μm through 1.49 μm taken by the
MASCS spectrometer during MESSENGER second Venus flyby on 5 June 2007. The analy-
sis included radiative transfer modelling to fit the observed spectra using vertical distribution
of aerosols and spectral properties of the UV absorber as free parameters. The imaginary part
of refractive index of the UV absorber was found to be centred at 0.34 ±0.03 μm with a full
width at half maximum of 0.14 ±0.01 μm assuming Gaussian shape of the absorption band,
thus suggesting the absorption a bit blue shifted with respect to the previous works. Figure 25
compares the UV-blue absorption derived from the MASCS/MESSENGER observations to
that of several species which spectral properties closely match the measurements. Several
candidates could account for the core absorption around 0.35 μm, but the main problem
is fitting the spectral slope at 0.4–0.5μm. The best agreement was found for an irradiated
version of S2O (Lo et al. 2003)andS
2O2or OSSO (Frandsen et al. 2016) if a single ab-
sorber is assumed. Other species including iron chloride have too narrow absorption to be
in agreement with the MASCS spectra. We note, however, that the analysis by Pérez-Hoyos
et al. (2018) did not include sulphur allotropes recently suggested by Carlson et al. (2016).
Recently, a possibility of microorganisms contributing to absorption of sunlight and con-
trasts in the clouds has been discussed by Limaye et al. (2018c). This hypothesis was based
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126 Page 34 of 61 D.V. Titov et al.
Fig. 25 Comparison of the
relative absorption in arbitrary
units (colour lines) of the
candidates for the near-UV
absorber proposed so far with the
absorption spectrum derived
from MASCS/MESSENGER
observations (grey area). The
black dashed line shows the best
fitting values (from Pérez-Hoyos
et al. 2018)
on the consideration that terrestrial-type biology may have evolved and could survive within
Venus clouds. The authors showed that the observed physical, chemical and spectral prop-
erties of the Venus cloud population can be matched by terrestrial microorganisms. Earth
clouds harbour many microorganisms, so a similar existence in the clouds of Venus cannot
be excluded.
In summary, identification of the UV-blue absorber (or absorbers) in the Venus atmo-
sphere is a problem far from being solved. Several single candidates provide good match to
the observed properties but none of them can fully and consistently explain the phenomenon,
including spectral properties, chemistry and dynamics. For instance, iron chloride (FeCl3)
whose presence in the Venus clouds agrees with chemical models (Krasnopolsky 2018)
does not seem to fit well the observed spectral behaviour (Fig. 25). On the other hand, the
abundance of disulphur dioxide (S2O2) that provides the best spectral fit deviates from the
profiles computed by Frandsen et al. (2016). Krasnopolsky (2018) also showed that the re-
quired S2O2abundance is in contradiction with state-of-the-art photochemical models. As
most of single candidates do not match the MASCS observations at 0.4–0.5μm(Fig.25)
Pérez-Hoyos et al. (2018) argued in favor of a second absorber that could complement the
species absorbing at shorter wavelength, such as S4, for instance. This echoes the tentative
conclusion by Titov et al. (2012) about likely presence of an additional species whose ab-
sorption extends to visible wavelengths to explain the cloud features seen by VMC/VEX
at 0.513 μm. Other constituents. There are hints of further non-sulphur cloud particle con-
stituents from several sources. As mentioned above, analysis of the cloud material acquired
by the aerosol collector/pyrolyser on the Vega probes by the mass-spectrometer suggested
presence of chlorine and sulphur. Inferred abundance in particulates in the cloud region
were 0.3 mg/m3for chlorine, compared to 2 mg/m3for sulphur (Surkov et al. 1986).
Even more intriguing are the results from Soviet X-ray radiometer experiments. Analysis of
aerosols on Venera-13 and -14 indicated presence of sulphur, chlorine and iron (Petryanov
et al. 1981), while similar experiments on Vega 1 and 2 found little iron but strong sig-
natures of sulphur, chlorine and phosphorus in the lower cloud (Fig. 26) (Andreichikov
1987). This could plausibly exist in the form of phosphoric acid droplets (H3PO4) with
P4O6as a possible gaseous species acting as a reservoir for phosphoric acid formation
(Krasnopolsky 1989). The sub-cloud hazes found at altitudes of 33–48 km at temperatures
too elevated for sulphuric acid droplets to exist, so might be explained by phosphoric acid
droplets. The lack of detection of any phosphorus-containing species from any other in-
vestigations makes it difficult to proceed further with evaluation of the possibility of the
phosphoric acid in the clouds, but it is clear that this is an intriguing area for future investi-
gation.
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Clouds and Hazes of Venus Page 35 of 61 126
Fig. 26 Accumulation of
chlorine, sulphur, and
phosphorus on the filter of
Vega-2 X-ray radiometer as a
function of altitude (from
Andreichikov 1987)
Finally, we should mention also the possibility of near-surface hazes. A recent reanaly-
sis of Venera-13, -14 descent probe spectrophotometer data found a sharp decrease of light
levels at 1–2 km altitude, interpreted as indicating a detached layer of aerosols of unknown
nature at this level (Grieger et al. 2004) The authors point out that its altitude is similar to
that at which radar-bright deposits attributed to metallic condensate on the mountain tops
have been found (e.g. Schaefer and Fegley 2004). The aerosol layer could also be asso-
ciated with volcanic ash or dust lifted by wind. Further investigation of such low-altitude
hazes may be possible by radar investigations or near-infrared spectroscopy on the night-
side.
6 Modelling Chemistry and Microphysics of the Clouds
The vertical structure and makeup of the system of Venus clouds and aerosols has been
described in the preceding sections. Here, we focus on modelling investigations of the pro-
cesses that operate within the clouds and aerosols of Venus. To facilitate that discussion, we
include a schematic figure that shows the interaction between the chemical, microphysical,
dynamical, and radiative processes that occur in the Venus clouds and lead to their long-
term stability, and their short-term variation (Fig. 27). The goal of microphysical modelling
of the Venus clouds is to characterize the physical properties of the aerosols in the Venus
atmosphere, through simulation and comparison with observations, so that we can better
elucidate these coupled drivers of cloud formation and dissipation. Hence, we will be able
to better leverage observations of changes in cloud morphology to measurements of atmo-
spheric circulation, chemistry, and radiative balance. However, the cloud system of Venus is
a highly coupled system, with many forcings, few of which are well-characterized.
6.1 Microphysical and Chemical Processes in the Clouds
Photochemical reactions involving water vapour and sulphur species shown schematically
as “H2O+SO2+sunlight” in Fig. 26, potentially including hundreds of possible reactions
(Krasnopolsky 2012; Zhang et al. 2012), produce sulphuric acid vapour in the middle at-
mosphere of Venus. Peak production is modelled to occur at about 62 km (Esposito et al.
1979), but even this is dependent upon the modelled pathway from vapour to cloud, since
direct measurement is very difficult because the concentrations of sulphuric acid number
density at these altitudes are so slight due to its very low vapour pressure.
At sufficient supersaturations, this sulphuric acid vapour will condense into cloud
droplets; this process can either occur homogeneously or heterogeneously (Seinfeld and
Pandis 1998). Homogeneous nucleation is conceptually simpler—the only requirement be-
ing sufficiently large concentrations of sulphuric acid molecules that collisions between
molecules are frequent enough that a cluster of nucleated acid molecules (sometimes called
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126 Page 36 of 61 D.V. Titov et al.
Fig. 27 A schematic figure of the Venus atmospheric clouds and hazes, provided as a guide to the discussion
in the text
a germ and having a radius on the order of nanometres), will remain intact long enough
for diffusional growth processes to begin to take over. This generally requires rather large
supersaturations. Heterogeneous activation of droplets, on the other hand, generally (but not
of necessity) requires a soluble condensation nucleus in order for the activated particle to
overcome the Kelvin barrier (below which radius, a very small droplet will evaporate more
quickly than it can grow, because the small radius of curvature increases the equilibrium
vapour pressure; the addition of a solute reduces this effect, allowing smaller radius par-
ticles to persist and then grow condensationally). This activation process will proceed at
lower supersaturation rates, but does require a condensation nucleus that is soluble in order
to do so. This is possible in the Venus clouds, especially in the lower and middle condensa-
tional clouds, because we have observational evidence for a population of submicron sized
particles below the main cloud deck that cannot be sulphuric acid, and likely are involatile
(Knollenberg and Hunten 1980). But since we do not have a firm understanding of the com-
position of the particles that can serve as potential cloud condensation nuclei in the Venus
atmosphere, simulations of this physical process necessarily involves the addition of a rela-
tively unconstrained knob in a microphysical model. Furthermore, the most likely candidates
for these cloud condensation nuclei are polyatomic sulphur species—which are necessary
products of the aforementioned photochemistry (e.g., Yung et al. 2009)andforwhichthere
is some observational evidence (e.g. Toon et al. 1984)—whichhavebeenshowntobein-
capable of serving as condensation nuclei, since they are completely insoluble in sulphuric
acid (Young 1983).
These photochemically produced sulphuric acid cloud droplets tend to form a several
kilometres thick cloud deck above about 57 km, which Fig. 26 labels the photochemical
cloud. These upper cloud droplets were determined to be spherical droplets with effective
radius of about 1 micron, composed of sulphuric acid by polarimetry measurements in the
1970s (Hansen and Hovenier 1974). Above this is a haze of particles that appears to be com-
posed of sulphuric acid and appears to be at least at times, bimodal (Kawabata et al. 1980;
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Clouds and Hazes of Venus Page 37 of 61 126
Wilquetetal.2012; Luginin et al. 2016). The upper haze extends to about 100 km. Above
the region of photochemical production the haze vertical structure is governed by dynamical
processes such as eddy diffusion and Stokes settling of the particles. Historically, the defini-
tion of this upper haze has varied, and this has led to some confusion and misunderstandings
in the community. In one sense it refers to the population of particles located above the main
cloud deck defined by the photochemistry and condensational processes occurring between
about 45 km and 65 km. At the same time, Kawabata et al. (1980) referred to haze particles
and cloud particles that existed in the top optical depth of the clouds accessible by polarime-
try. In this description, the micron sized “cloud” particles and the submicron sized “haze”
particles sometimes co-existed in the uppermost optical depth, and sometimes the popula-
tion was dominated by the submicron sized particles. The point of this tangent is to draw the
reader’s attention to the inconsistencies in the definitions in use in the literature, and to take
the time to state here that in this paper, “upper haze” refers not to a size of particle but rather
to the population of particles in the topmost optical depth that can be sensed by polarimetry.
The characteristics of the upper haze are discussed in Sects. 3and 4above. When two
modes of particles are seen in the upper clouds and hazes, they tend to have typical radii
of around one micron—consistent with the particles first characterized by Hansen and Hov-
enier (1974)—and a submicron mode. At about 57 km, a transition region is observed in
most descent probe data (Knollenberg and Hunten 1980). Perhaps coincidentally, this alti-
tude is consistent with the altitude below which the ambient temperature warms to a point
at which frozen particles of sulphuric acid can no longer be supported (see Sect. 8.1). Also
at this altitude, the tops of the middle clouds can cool sufficiently during the Venusian night
to drive convection in the middle and lower cloud decks (Baker et al. 2000; McGouldrick
and Toon 2008; Imamura et al. 2014). Furthermore, this vertical mixing, driven both by this
cooling at 57 km and heating at cloud base due to absorption of upwelling IR, draws large
quantities of sulphuric acid vapour upward in to the cloud from a reservoir below. Thus, in
Fig. 27 the cloud deck below 57 km is called the condensational cloud that corresponds to
the middle and lower clouds in the earlier terminology. Particle sizes in this region of the
cloud tend to be somewhat larger than those in the upper clouds and hazes. A possibly tri-
modal distribution of cloud particles was observed in the condensational cloud by the Large
Probe Cloud Particle Size Spectrometer (LCPS) of Pioneer Venus (Knollenberg and Hunten
1980). Furthermore, the typical radius of the 1-micron size mode of particles was observed
to increase with depth, both in the aforementioned in situ measurements, as well as in sub-
sequent remote sensing analysis of the near-infrared spectral “windows” by Carlson et al.
(1993) and Grinspoon et al. (1993).
Total opacity variations in this condensational cloud are largely responsible for the
inhomogeneities seen in the emitted near infrared (Crisp et al. 1989). Descent probes
from Venera and Pioneer Venus indicated large spatial variations in the vertical profile
of mass loading of these clouds, especially in the lower cloud between about 48–50 km
(Ragent and Blamont 1980; Marov et al. 1980). And analysis of Galileo NIMS obser-
vations of the NIR spectral “windows” that had been discovered by Allen and Crawford
(1984) suggest that these variations can be attributed to both changes in total mass of
the clouds as well as changes in the particle size distributions (Grinspoon et al. 1993;
Carlson et al. 1993). Since the vapour pressure of sulphuric acid is so low and so sen-
sitive to temperature and concentration, subtle changes in microphysical properties, or in
temperature or in local winds can all potentially drive these changes in cloud characteristics
(McGouldrick and Toon 2007).
At altitudes below about 40 km, the sulphuric acid vapour is thermally decomposed, ul-
timately back into SO2and H2O (and likely other sulphur, oxygen, and hydrogen bearing
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126 Page 38 of 61 D.V. Titov et al.
species (Krasnopolsky 2013). Although this is many kilometres below the clouds, it is a pro-
cess that is highly dependent upon temperature and local stoichiometric composition. Thus,
under the right conditions, this decomposition can have a feedback on the physical proper-
ties of the cloud particles above 48 km. Thus, in order to fully understand the condensational
clouds of Venus, we need to understand the chemical drivers, as well as microphysical; just
as we do for the photochemical cloud.
Alternate compositions of the Venus clouds are possible at some altitudes (see previ-
ous section), including possibilities for rather exotic possibilities in the deep atmosphere of
Venus. However, to date, nearly all in situ and remote sensing measurements (with the most
notable exception being the largest size mode of particles observed by LCPS) are consistent
with sulphuric acid having a weight percent between 75% and 90%, so this remains the pre-
ferred identification of the composition of both the photochemical and condensational cloud
decks of Venus.
6.2 Modelling the Venus Clouds
Modelling of aerosol microphysical properties and processes involves the simultaneous so-
lution of several differential equations that govern the growth and transport of the aerosol
particles. Namely, nucleation of particles (P), condensational growth and evaporation of the
aerosols (G) must be considered including total evaporation (L), as well as coagulation (and
coalescence) between particles (integrals involving Kcoag), and transport by diffusion (Kdiff),
winds (w) and gravitational sedimentation (vfall):
∂N(m,z,t)
∂t +∂
∂zN(m,z,t)w(z) −vfall(m,z,t)
−∂
∂zρ(z)Kdiff(z) ∂
∂zN(m,z,t)
ρ(z) 
=P(m,z,t)−L(m, z, t )N (m, z, t)
+m
0
Kcoag(m1,m−m1)N(m1,z,t)N(m−m1,z,t)dm
1
−N(m,z,t)∞
0
Kcoag(m1, m)N (m1,z,t)dm
1−∂
∂mN (m, z, t )G(m, z, t)
Here, the particle concentrations (N) are functions of mass (M), altitude (z)—or, more
generally, spatial location (x,y,z)—and time (t).
Thus, the growth of individual aerosol particles and the evolution of aerosol populations
are dependent upon temperature, the concentrations of the condensing species, the concen-
trations of solutes within the aerosols, the diffusional processes of the dry atmosphere, and
the number density of the particles themselves (more correctly, the size distributions). In
addition, transport processes can change the environment in which the particles exist, hence
changing their growth rates. Three-dimensional winds may advect particles and gasses and
thermal energy, while the particles themselves, if large enough, may fall relative to the am-
bient air. And, as discussed above, varying aerosol particle size distributions and trace at-
mospheric species can alter the local radiative properties, leading to significant changes in
heating and cooling rates. Since the system can be so heavily coupled, previous researchers
have endeavoured to understand the cloud system by making various assumptions about one
or more of the atmospheric or microphysical properties in order to reduce the complexity
of the problem to a manageable level. Next, we shall describe some of the microphysical
modelling efforts that have transpired since the publication of the Venus II book in 1997.
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Clouds and Hazes of Venus Page 39 of 61 126
Fig. 28 Size distribution of
Venus aerosols (both volatile and
involatile) at 50 km for several of
the simulations by James et al.
(1997). This figure shows the
emergence of a third mode in the
size distribution when solute
effects are taken into account
(increasing the efficiency of
condensational growth by vapor
diffusion), and when the
efficiency of condensational
growth is reduced artificially by
altering the diffusional flux of
molecules to an existing droplet
James et al. (1997) developed a one-dimensional model of the condensational clouds of
Venus in which the upper cloud was a boundary condition flux of particles into the model
domain, determined from in situ measurements of the downward mass flux of upper cloud
particles at 57 km. They also stipulated a profile of eddy diffusion coefficient that was con-
stant with time, and that was based on existing understanding of convective vertical motions
in the middle and lower cloud decks. They were able to reproduce the vertical profiles mass
loading of the condensational cloud (as compared with in situ measurements by Pioneer
Venus descent probes). However, they could only reproduce the trimodal structure of the
size distribution in the lower clouds by invoking solubility in the activation of cloud particles
in the condensational cloud (Fig. 28). As mentioned above, however, since the composition
of the involatile particles in the lowest part of the Venus cloud system is unknown, their
modelled solubility properties were necessarily arbitrary.
Imamura and Hashimoto (1998) sought to understand connections between large-scale
dynamical transport and observable aerosol properties and distribution. They explored the
consequences of a putative near-global Hadley-like meridional circulation on the latitudi-
nal distribution of aerosol opacity. Unlike James et al. (1997), they included photochemical
production of sulphuric acid vapours in order to adequately model the formation of the up-
per photochemical clouds as well as the deeper condensational clouds. However, they did
not attempt to simulate evolution in the size distributions of the particles, instead treating
the particles as comprised of two monodisperse populations, having radii of 1.15 μmand
3.65 μm, consistent with the in situ observations of modes 2 and 3. They found that their
combination of simplified chemistry, meridional transport, and simplified aerosol micro-
physics was able to largely reproduce the observed latitudinal distribution of aerosol mass
loading (Fig. 29), aerosol vapour mixing ratios, and average radiance that is seen in the
nightside NIR emission. They perceived a build-up of mass in the upper cloud at high lati-
tude, arising from the accumulation of photochemical cloud particles as the sulphuric acid
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126 Page 40 of 61 D.V. Titov et al.
Fig. 29 The mass loading of the sulphuric acid water solution droplets in the nominal model of Imamura and
Hashimoto (1998) which shows condensational cloud mass enhancements at equatorial latitudes, and upper
cloud mass enhancements at polar latitudes
is exposed to sunlight during its trek from equatorial latitudes toward the poles. They also
found an enhancement of cloud thickness near the equator, which they attributed to the up-
welling branch of the Venusian meridional circulation. However, since their simulations did
not account for changes in microphysical properties (such as fall velocities, growth rates,
and coalescence rates) that are functions of particle size, the application of their results to
more localized events is not yet possible.
Imamura and Hashimoto (2001) later developed a bin-resolved microphysical model,
similar to that of James et al. (1997), that they used to demonstrate that the trimodal size
distribution observed in the condensational cloud can be obtained by mixing a column ex-
periencing upwelling with another experiencing downwelling, and allowing the merged col-
umn to evolve with time (Fig. 30). This simulation showed that the size distribution existing
in the condensational cloud at any time is as much dependent upon the current conditions as
it is on the history of local eddy activity.
McGouldrick and Toon (2007) reproduced the work of James et al. (1997) with an up-
dated version of the same model; but then went on to demonstrate the role of the radiative-
dynamical feedback in the evolution and sustenance of the Venusian condensational clouds.
Rather than employ an eddy diffusion profile that was constant in time as had been done pre-
viously, they calculated an eddy diffusion coefficient that was based on the static stability
of the local atmosphere, parameterized according to the Richardson number, which repre-
sents the ratio of Brunt-Vaisala frequency to the vertical shear of the horizontal winds. Their
model incorporating this radiative dynamical feedback was able to reproduce the aerosol
mass loading and backscatter cross section observed by the Pioneer Venus LCPS (Fig. 31).
They also showed that waves propagating through the clouds, and vertical winds consistent
with those observed by the in situ Vega balloons were capable of producing the observed
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Clouds and Hazes of Venus Page 41 of 61 126
Fig. 30 The evolution of the size distribution of cloud particles simulated by the model of Imamura and
Hashimoto (2001), at three altitudes over a time span of ten hours after the initial mixing of a column that
had experienced upwelling with another that had experienced downwelling
Fig. 31 Vertical profiles of the extinction coefficient and the backscattering cross section at 0.56 μm calcu-
lated from the simulations of McGouldrick and Toon (2007). The solid line represents the results using a static
eddy diffusion profile; the broken lines are the results from simulations incorporating the radiative-dynamical
feedback
clearings in the clouds in timescales consistent with the observed changes. McGouldrick
and Toon (2008) simulated small scale convective overturning by applying a very simple
convectively overturning wind field that was not responsive to the local temperature struc-
ture. The simulation was valid only a for a few hours of modelled time, but showed that
while convective cells were capable of generating NIR nightside emitted radiance contrasts
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126 Page 42 of 61 D.V. Titov et al.
comparable to those observed between clouds and holes in timescales of as small as hours,
they also demonstrated that the radiative-dynamical feedback applied in McGouldrick and
Toon (2007) needed to be employed in three-dimensional models in order to adequately
simulate the evolution of the condensational clouds of Venus on longer timescales. In each
of these simulations, McGouldrick and Toon included the upper clouds (and hazes) as a
boundary condition to their models, since they were focused primarily on the microphysics
of the condensational clouds. Thus any potential feedbacks caused by interaction between
the photochemical and condensational clouds were absent.
Vertical transport of H2OandH
2SO4in both gas and condensed phases was calculated
by Krasnopolsky (2011,2015) for the conditions of the main clouds (45–75 km) and the
haze (70–110 km), respectively. The model explains the observed reduction of the H2O
mixing ratio from the cloud bottom to the top by an order of magnitude and formation
of a layer of gaseous sulfuric acid centered at the lower cloud boundary. The models were
calculated for low, middle latitudes and 60◦. Formation of the lower cloud layer is stimulated
by a strong vertical gradient of gaseous sulfuric acid near the boundary (Krasnopolsky and
Pollack 1994). Concentration of sulfuric acid decreases from 98% at the cloud bottom near
48 km to 80% in the upper cloud layer, then reaches a minimum of 70% at 90 km and
increases to ≈100% at 100–110 km.
Observations of occultation data by SPICAV on Venus Express indicated an “inversion”
of the SO2profile whereby significantly larger concentrations of the gas were observed than
would be expected if the concentration of SO2decreased with altitude as a consequence of
diminishing atmospheric density and photochemical destruction of SO2. Zhang et al. (2012)
suggested that the vertical transport of sulphuric acid aerosols to altitudes >90 km where
they would evaporate and then the resultant sulphuric acid vapour could be photolyzed to
produce the observed SO2enhancement. Others have shown that the role of available water
in the vibrational overtone initiated chemistry that would drive this photolysis of sulphuric
acid is a significant factor in its ultimate atmospheric fate (Vaida and Donaldson 2014). This
demonstrated the need to consider both microphysical processes as well as chemistry in
order to adequately simulate the behaviour of the upper clouds and hazes of Venus.
Gao et al. (2014) and Parkinson et al. (2015) found a bimodal size distribution of aerosols
in the upper hazes of Venus and a trimodal distribution in the lower cloud by considering
activation of sulphuric acid droplets onto meteoritic dust and a monodisperse distribution of
polyatomic sulphur species that are products of the same photochemical process (Fig. 32).
However, this latter microphysical behaviour is inconsistent with the physical properties
of polysulphur and sulphuric acid (Young 1973). Gao et al. (2014) and Parkinson et al.
(2015) see a correlation between the production rate of polysulphur and the presence of
larger “rain” droplets of sulphuric acid deeper in the cloud. When the typical polysulphur
particle produced is small (and numerous), oscillations in the size distribution of the deeper
sulphuric acid droplets that occasionally produce rain; they do not see these oscillations
when the production rate is smaller. Parkinson et al. (2015), using a very similar model to
Gao et al. (2014), finds that this effect only occurs equatorward of about 60 degrees latitude.
However, unlike James et al. (1997) and McGouldrick and Toon (2007,2008), Gao et al.
(2014) and Parkinson et al. (2015) fix the vertical profile of water vapour in their models.
Imamura et al. (2014) showed that, perhaps counter-intuitively, a consequence of vary-
ing insolation due to the atmospheric super-rotation was to enhance vertical mixing (i.e.,
unforced convection) during the night relative to the day. Since the peak of absorption of
the incident solar radiance occurs in the vicinity of the cloud tops, and since the clouds
and atmosphere are so optically thick in the infrared that the radiative cooling to space is
relatively insignificant at most infrared wavelengths below an altitude of about 57–60 km
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Clouds and Hazes of Venus Page 43 of 61 126
Fig. 32 The size distribution as
a function of altitude from the
nominal model of Gao et al.
(2014). Note the existence of a
trimodal size distribution in the
deep cloud, and suggestions of a
bimodal cloud structure above
60 km
Fig. 33 Time series of
root-mean-square vertical
velocity and convective heat flux
from the dynamical models of
Imamura et al. (2014), showing
the significant reduction in
convection to be expected at
midday in lower latitudes in the
Venus atmosphere and the more
consistent levels of convective
flux that occur at higher latitudes
(near the upper cloud to middle cloud transition), the primary effect of insolation is to sup-
press cooling from the top of the condensational cloud deck, hence suppressing the build-up
of the thermal gradient necessary to drive free convection in the middle clouds (Fig. 33).
Since the altitude range of the upper clouds is so stable to overturning, the variations in
insolation there are still too small to drive convection above 60 km. Imamura et al. (2014)
did not directly consider the effects of these variations in vertical mixing would have on
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126 Page 44 of 61 D.V. Titov et al.
the cloud structure by coupling their dynamical model to a microphysical cloud model; but
previous work by McGouldrick and Toon (2007,2008) suggested that these subtle changes
in the radiative-dynamical feedback could produce small but significant effects on the cloud
particle distribution.
To date, very little modelling of the Venusian clouds has considered compositions other
than sulphuric acid. However, as mentioned in the previous section, there is observational
evidence that the aerosol could include additional compositions. Chlorine was detected in
the aerosols by the Vega probe. Recent work by Krasnopolsky (2017) suggests that iron
chlorides could represent a significant fraction of the mode 1 aerosol mass, and could be
responsible for the near ultraviolet absorption. Further supporting this hypothesis is the re-
cent work by Sandor and Clancy (2012) who found that concentrations of hydrochloric acid
above the cloud tops are somewhat less than would be expected from current photochemical
models. Furthermore, Sandor and Clancy (2017) find that these HCl concentrations do not
vary diurnally. One possibility is that Vega collected at least some sulphuric acid droplets
that also contained ferric chloride.
7 Radiative Effects of the Clouds
Clouds are one of the main factors forming the Venus climate due to complete coverage of
the planet, large total opacity and specific optical properties (Crisp and Titov 1997; Titov
et al. 2007,2013). Sulfuric acid aerosols almost conservatively scatter the solar light in the
range from UV to near-IR (∼2.5μm). As a result, more than 75% of the incoming solar
flux is returned back to space and, on average, only ∼3% of the solar flux incident at the top
of the atmosphere reaches the surface. This results in a surprising fact: even though Venus
is 30% closer to the sun, it receives less solar energy than the Earth due to the bright cloud
“blanket” shrouding the planet. On the other hand, strong absorption at longer wavelengths
makes sulfuric acid clouds a powerful greenhouse agent responsible for trapping thermal
radiation at λ>2.5μm in the lower atmosphere. The unknown near-UV absorbers plays
significant role in deposition of the solar energy due to that its absorption band is located in
the UV-blue range in which the solar radiation is the most intensive. Figure 34 shows the
mean spectrum of Venus as seen from space. It has three components: the reflected solar
light dominating at λ<4μm, thermal radiation emitted to space by the cloud tops and the
mesosphere that prevails at λ>4μm, and emissions leaking from the lower atmosphere
through the so-called spectral transparency “windows” in the near-IR (0.6–2.5μm). This
section reviews the effects that Venus clouds have on the radiative energy budget and com-
plements more general and comprehensive discussion of this topic by Limaye et al. (2018a).
Inhomogeneous distribution of the unknown near-UV absorbers at the cloud tops that
produces famous contrast markings (Fig. 13) strongly affects the radiative energy deposition.
The observed global latitudinal variations of the UV-blue albedo imply that the UV-dark
tropics receive by a factor of 3–4 greater amount of solar energy than the UV-bright high
latitudes. This contrast in energy deposition between the low and high latitudes is further
enhanced by gradual equator-to-pole decrease of sun elevation angle. Latitudinal, spatial
and temporal changes of the planetary albedo affect the radiative energy deposition pattern
and thus have an impact on the atmospheric dynamics. The unknown UV-blue absorber was
found to be confined to the upper cloud (57–65 km) (Ekonomov et al. 1984). This results in
absorption of about half of radiative energy that the planet receives from the sun at the cloud
tops, thus creating energy deposition pattern remarkably different from the other terrestrial
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Clouds and Hazes of Venus Page 45 of 61 126
Fig. 34 General view of the
Venus spectrum at a resolving
power λ/λ ∼200 as seen from
space: the reflected solar light
(dashed line), the thermal
emission from the cloud tops and
mesosphere (dotted line), and the
night side emission escaping
from the lower atmosphere (solid
line). Positions of strong gaseous
and aerosol absorption bands are
marked (from Titov et al. 2007)
planets (Crisp and Titov 1997). This drives an unusual global circulation in the form of
retrograde zonal superrotation.
Venus returns energy to space mainly through thermal emission from the cloud tops in
the range 10–50 μm(Fig.34). The outgoing thermal flux depends on the temperature and
aerosol distribution at the cloud top level. Recent progress in the observations of the temper-
ature and cloud vertical structure and their variability over the planet enabled re-assessment
of the radiative energy balance in the mesosphere including the cloud tops. Lee et al. (2014)
performed detailed numerical analysis of the sensitivity of the outgoing thermal radiation to
temperature and cloud structure in the Venus mesosphere using recent observation results
(see also Limaye et al. 2018a, this Book and Sect. 2above).
Sensitivity of the outgoing thermal flux on the cloud top altitude decreases with latitude
due to decreasing temperature gradient in the mesosphere. Small aerosol scale height in
polar regions (Fig. 11) is responsible for a pronounced cooling rate peak at the cloud top.
This feature corresponds to sharp upper boundary of the cloud, making the cloud top an
effective emitter of thermal radiation and providing a positive feedback between the temper-
ature and cloud structures. The cooling rate value increases from 5 K/day to ∼30 K/day
as the aerosol scale height decreases from 6 km to 1 km regardless of thermal structure.
Effective cooling at the sharp cloud top can be considered as one of the factors maintaining
temperature inversion in the “cold collar” regions that, in turn, creates convectively stable
atmosphere with suppressed vertical mixing.
Solar heating also depends on the cloud structure and, first of all, on the presence of the
unknown near-UV absorbers which contribution to the overall heating is 30–60%. At low
latitudes the solar heating at the cloud top is more than twice as strong as thermal cooling,
meaning that at noon the solar heating at the cloud top dominates. In the “cold collar” re-
gion the solar heating and thermal cooling rates at the cloud top are comparable. In the polar
regions thermal cooling prevails. Figure 35 shows the net radiative energy budget in the
Venus mesosphere (Lee et al. 2014). Above the cloud top, the net cooling prevails poleward
from ∼50◦. At lower latitudes the mesosphere experiences slight net heating. This result is
in agreement with the earlier studies in that the Venus mesosphere is in radiative disequi-
librium and that meridional overturning of the atmosphere above the clouds is required to
maintain the observed temperature field.
Haus et al. (2016) performed the most complete and comprehensive study of the radiative
energy balance in the Venus atmosphere using state-of-the-art radiative transfer technique
and improved atmospheric model based on the Venus Express observations. The clouds
were found to strongly affect both atmospheric cooling in thermal-IR range and heating at
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126 Page 46 of 61 D.V. Titov et al.
Fig. 35 Latitude–altitude field of the net radiative forcing in the Venus mesosphere (Lee et al. 2014)
UV to near-IR wavelengths mainly due to changes in cloud composition and distribution
of the UV-blue absorbers with latitude. Cooling in the upper cloud and haze (65–80 km)
increases with mode 2 particles abundance, while mode 3 enforces cooling at 55–65 km.
Solar heating rates strongly decrease with latitude due to decreasing Sun elevation angle
while the effect of cloud structure is less significant that for cooling rates. The near-UV
absorbers are responsible for up to 10 K/day additional heating at the cloud top. Its presence
results in strong net heating in the upper cloud (60–70 km) during the day especially in low
and middle latitudes vanishing towards the poles. Below about 53 km an excessive heating
by the clouds results in overall very weak net heating at all latitudes. Comparison of radiative
energy calculations by Lee et al. (2015) and Haus et al. (2016) at 60–80 km indicate that
the results are very sensitive to the cloud model as well as latitudinal variations of the cloud
structure and distribution of the near-UV absorbers.
Thermal radiation leaking from the hot lower atmosphere and the surface through the
transparency “windows” gives minor contribution to the global radiative energy balance
due to weakness of the emission escaping through the thick cloud layer (Fig. 34). However
the radiative effect of the thermal emission from the hot lower atmosphere the deep cloud
could be much stronger. Measurements of the flux in the near-IR transparency “windows”
from orbit give an important clue on distribution of opacity over the globe (Figs. 12,20).
They indicate a factor of 5–7 decrease of the flux between tropics and polar regions that
corresponds to about a factor of 2 increase in the cloud opacity.
8 Lightning
Conventional wisdom suggests that lightning should not occur in the atmosphere of Venus.
However, the processes that lead to the build-up of charge and subsequent electric field
breakdown in terrestrial lightning is still poorly known and is an emerging field of active re-
search (see for example the recent joint ISSI-Europlanet workshop on Planetary Electricity
(Bonnet and Blanc 2008)). These advances, along with more recent observational programs,
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Clouds and Hazes of Venus Page 47 of 61 126
have re-opened the case for lightning on Venus. We are only beginning to understand the
requirements necessary for lightning on Earth, what does the application of that knowledge
to Venus tell us? In order to trigger a lightning discharge, it is necessary to build up a suffi-
cient quantity of charge that is separated by a distance of a few km. Once the electric field
produced by this charge build-up and separation exceeds the breakdown potential, a series
of discharges occurs at a frequency of thousands per second (LF) to millions per second
(HF), which the human eye perceives as a single lightning stroke (Zarka et al. 2008). Sec-
tion 8.1 describes the evidence for and against the existence of potential mechanisms that can
drive lightning processes in the Venus atmosphere. Nevertheless, as noted by Stolzenburg
and Marshall (2008) in the aforementioned review on Planetary Electricity, “As is evident
from this review, the electrical nature of thunderstorms on Earth is neither simple nor fully
known.”
The search for lightning on Venus has proceeded haltingly, with many reliable detections,
and many reliable non-detections, across a range of observational techniques. A recent sum-
mary by Takahashi et al. (2008) guides the summary of past observations in Sect. 8.2,but
a forthcoming review by Lorenz (2018) promises to pursue these observations and their in-
tercomparisons in a far more comprehensive way than that of which we are capable in this
more broadly focused review.
8.1 Observational Evidence for and Against Venusian Lightning
The first observational evidence of lightning on Venus was from the Venera 9/10 scan-
ning spectrometer (Krasnopolsky 1983) that observed random flashing over a region of
5×104km2. The mean flash duration was 0.25 s, flash optical energy 2 ×107J, and a rate
of 45 km−2yr−1extrapolated over the whole planet. However, only one collection of events
was observed during that mission. Subsequently, Venera-11, -12 detected radio emissions
consistent with lightning discharges (Ksanfomaliti 1980). These tended to be at somewhat
lower frequency than typical terrestrial discharges, however (Takahashi et al. 2008). The
Pioneer Venus Orbiter Electric Field Detector (OEFD) also detected very low frequency
signals (Scarf et al. 1980), but these signals were open to multiple interpretations.
The Galileo and Cassini spacecraft carried instrumentation designed to measure and
characterize lightning on Jupiter and Saturn, respectively (Gurnett et al. 1991). During their
Venus flybys, Galileo detected six pulses with radio frequency greater than 1 MHz, but the
Cassini spacecraft detected nothing in the range 0.125 MHz to 16 MHz, despite readily de-
tecting signals of lightning during its Earth flyby and with flyby distances much closer to
Venus than to Earth during each respective flyby. These discrepancies suggest that if light-
ning exists at Venus, then either Venusian lightning is much weaker than terrestrial in inten-
sity (hence could not be detected), is a sporadic phenomenon (that Galileo was lucky enough
to encounter, but Cassini was not), or occurs with much lower stroke frequencies during a
discharge (hence detectable only in LF ranges (Zarka et al. 2008), which Cassini may not
have been able to detect). Most recently, the Venus Express magnetometer has reported
detection of low frequency radio detections of lightning at Venus. Circularly polarized elec-
tromagnetic waves with frequencies around 100 Hz have been interpreted by Russell et al.
(2007) as whistler modes that would indicate a definitive detection of lightning in the atmo-
sphere of Venus with frequency of occurrence being similar to that of Earth. More recent
analysis suggests that both of these processes are likely occurring. Russell et al. (2013) con-
cludes that linearly polarized signals detected by the Venus Express magnetometer at very
low frequencies (less than about 20 Hz) arise from interactions at the solar wind interface
and cannot be attributed to lightning. However, they also conclude that the circularly polar-
ized signals detected by Venus Express with slightly higher frequencies can only have an
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126 Page 48 of 61 D.V. Titov et al.
atmospheric source and remain consistent with an interpretation as whistler mode radiation
having been generated by lightning in the atmosphere—likely the clouds—of Venus. It is
difficult to calculate a lightning occurrence rate from the Venus Express data, because it
could detect these signals only when within the ionosphere, low above the Northern hemi-
sphere of the planet; however, an approximate estimate suggests that the lightning flash
occurrence rate inferred from Whistler wave detections is of a similar order of magnitude to
that found on Earth (Russell et al. 2007).
Optical searches for lightning—after the Venera- 9 ,-10 observation mentioned above—
also include both positive and negative detections. The Pioneer Venus star tracker was uti-
lized in a search for optical signatures of lightning in the Venus clouds, but was unsuccess-
ful (Borucki et al. 1991). However, with only about 83 s of observing time over a two year
search, this campaign was able to provide only an upper limit for lightning discharge fre-
quency. Hansell et al. (1995) performed a ground-based optical search for lightning which
resulted in 7 positive detections from about four hours of observation over eight nights, fur-
ther suggesting either the low frequency of occurrence or the weak signal (or both). The only
pre-planned and directed in situ optical search for lightning in the Venus clouds was carried
out by the Soviet Vega balloons, which carried optical detectors pointed downward while the
balloon floated at an equilibrium altitude of about 53 km. In principle lightning occurring at
altitudes near 57 km could have been detected by this search if not the strong background
due to surface emission in 1 μm transparency “window”. Finally, Cardesín-Moinelo et al.
(2016) conducted a search for lightning in the atmosphere of Venus in VIRTIS-M-Vis data
between 0.3μmand1.1μm. They found a number of flash events in the data set; but none
could be explained definitively and exclusively as having been produced by lightning. As
a consequence, they report that the lightning rates reported by Russell et al. (2011) cannot
be considered to be a global phenomenon. The Akatsuki orbiter, which arrived at Venus
in 2016, carries a Lightning and Airglow camera which is designed to search for lightning
in the 0.7774 μm OI emission band (Takahashi et al. 2008); its initial results have yielded
no lightning detections, but this still on the basis of a very small number of observations
(Takahashi et al. 2018).
There have also been chemical arguments for existence of lightning. Krasnopolsky
(2006) argued that if lightning would be the only source of nitric oxide in the lower at-
mosphere of Venus the observed NO mixing ratio of 5.5±1.5 ppb requires the lightning
energy deposition of 0.19 ±0.06 ergcm−2s−1that is even higher than that on the Earth.
8.2 The Microphysics of Venusian Lightning
For a complete discussion of the microphysics of the Venusian clouds, please see Sect. 6.
Here we discuss only the processes relevant to the production of putative lightning dis-
charges in the Venus clouds. Numerous charging mechanisms have been identified to have
varying levels of importance on Earth and we discuss those which seem to be relevant to
the Venusian atmosphere. The simplest case is ion-induced charging, which may be more
important on Venus than Earth, due to the lack of an intrinsic magnetic field, which would
allow greater numbers of cosmic rays and solar energetic particles to reach deeply into the
atmosphere, possibly to the altitudes inhabited by the clouds (Nordheim et al. 2015). Ion-
induced charging is a function of supersaturation and temperature (Aplin 2013), suggesting
that the lower clouds are the ideal location for this type of charging process to occur. Fur-
thermore, the lower vapour pressure of sulphuric acid and its larger permittivity than that
of water mean that smaller supersaturations will be required to drive ion-induced charging
in the Venus atmosphere. While Aplin (2013) suggests that the conditions for ion-induced
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Clouds and Hazes of Venus Page 49 of 61 126
Fig. 36 The sulphuric acid freezing/melting phase diagram. The solid line represents the freezing tempera-
ture of sulfuric acid as a function of weight percent (by mass), from Gable et al. (1950). Sulfuric acid above
the line is liquid, below is either frozen or supercooled liquid. Note that, in this diagram, temperature in-
creases vertically, so altitude increases downwards. Finally, also plotted, with a dotted line, is the “cooling
curve” of a typical Venus sulphuric acid aerosol, assuming temperature profile from the Venus International
Reference Atmosphere (Kliore et al. 1986), and water vapor mixing ratio vertical profile from Fig. 16b of
Pollack et al. (1993); individual Venus altitudes are indicated by symbols (+). The weight percents of the
monohydrate (SAM) and dihydrate (SAD) are indicated by the dash-dot vertical lines. For comparison, typ-
ical Earth Stratospheric conditions are indicated by the diagonal hatching on the lower left of the diagram.
A similar version of this figure appeared in McGouldrick et al. (2011)
charging are best in the vicinity of the lower clouds (due to the higher temperatures), super-
saturations can become very large indeed in the upper clouds and hazes (Gao et al. 2014;
McGouldrick 2017). It may be possible that these large supersaturations may overcome the
temperature effect to encourage ion-induced charging in the upper clouds of Venus.
Perhaps the most likely means of charge build-up on Venus aerosols would be tribo-
electric charging. On Earth, this usually requires co-existence of frozen and liquid phases
of water; but can also proceed due to interactions of a compositionally homogeneous pop-
ulation with a range of sizes (e.g. triboelectric charging of volcanic dust during an erup-
tion). On Venus, the solution of sulphuric acid and water that is the primary composition
of the cloud particles has a freezing point that is a function of both the temperature and the
acid mass concentration (which itself is a function of the water vapour concentration and
the temperature). Note in Fig. 36 that for typical Venus conditions (as estimated from the
VIRA equatorial model), the phase change occurs at 57 km, around the level of the tran-
sition between the upper (photochemical) and middle and lower (condensational) clouds.
Thus, interactions between frozen and melted sulphuric acid cloud droplets near 57 km, fol-
lowed by subsequent separation (whether gravitationally or via vertical shear of the zonal
wind), could lead to build-up of sufficient charge to lead to lightning discharges on Venus.
A complication is that in the laboratory, and in the terrestrial stratosphere, it has proven
exceedingly difficult to freeze sulphuric acid (Zhang et al. 1993). On the other hand, su-
percooled sulphuric acid (and in the Venus atmosphere, SAM (see Fig. 36)mayexistin
conditions 50 K below its freezing point) is known to become highly viscous and glasslike
(Horn and Sully 1999). In this case, collisions between cloud particles would behave more
like homogeneous triboelectric charging, in which differences in particle radius drive the
exchange of electrons (Forward et al. 2009). The smooth surfaces of the glasslike sulphuric
acid would encourage charge separation, and the magnitude of this charge separation would
be a function of the differences in the sizes of the colliding particles (Yair et al. 2008;
Lacks and Levandovsky 2007). Thus, if triboelectric charging of sulphuric acid cloud par-
ticles is the source of the putative Venusian lightning, then the mostly likely place for the
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126 Page 50 of 61 D.V. Titov et al.
charge build-up and separation to be occurring would be the region between the upper and
middle clouds.
A final possible mechanism for Venusian lightning is near-surface triboelectric charging
due to collisions of sediments near the surface lofted either by Aeolian or Volcanic pro-
cesses. This mechanism, however, has some challenges to overcome. For example, there is
in the continued lack of direct evidence for volcanism near the surface. The population of
particles suitable for saltation and the magnitudes of the near surface winds that would drive
such transport are poorly known. Finally, the large dielectric constant of such a dense atmo-
sphere suggests that very large charges must be accumulated on particles in order to achieve
a breakdown of the electric field.
9 Synthesis of the Cloud Properties
Vertical Structure Venus possesses the largest and the most complex cloud system among
the terrestrial planets. The clouds that completely shroud the planet are located in the alti-
tude range of approximately 50–70 km with upper and lower haze extending up to about
100 km and 30 km correspondingly, thus occupying the upper troposphere and the entire
mesosphere. The total opacity of the cloud at visible and near-IR wavelengths varies from
20 to 40 showing remarkable spatial and latitudinal variability. The vertical structure of
the cloud was established by Venera, Pioneer-Venus and Vega descent probes with several
layers identified from the measurements of extinction coefficient and number density (see
Esposito et al. 1983,1997 and references therein). Venus Express observations significantly
complemented and detailed the earlier picture especially concerning the properties of the
mesospheric haze as well as spatial and latitudinal variability of the cloud properties.
Figure 37 shows a synthesis of the latitude-altitude cloud structure and microphysical
properties. The upper haze occupies almost the entire mesosphere from about 100 km down
to the cloud top (blue line in Fig. 37), defined as the altitude of vertical opacity τ=1at
visible wavelengths. This level is located at ∼70 km at low latitudes and gradually descends
to ∼62 km at the poles (Figs. 12,37) thus changing by about two atmospheric scale heights
and creating a vast polar depression The aerosol scale height in the upper cloud follows a
similar trend decreasing from 4–5 km at low-middle latitudes to 1–2 km in the “cold collar”
(60–70◦) and polar regions correlating with the mesospheric temperature field (Fig. 37).
Remarkably, the “cold collar” also marks the latitudinal boundary at which the total cloud
opacity at 1 μm abruptly increases by a factor of 2–3 (Fig. 20).
The green line in Fig. 37 marks the tropopause (Tellmann et al. 2009). This boundary
separates the region of photochemical production of sulphuric acid (upper cloud and haze)
from the condensational cloud below (see Fig. 27) and can be considered as the physical
base of the upper cloud. The middle and lower clouds are located below the tropopause
where convection dominates. The cloud base is located at 48–50 km at low and middle lat-
itudes. There are evidences from the sounding in the near-IR transparency “windows” and
radio-occultations that in the polar regions the cloud deck extends to even below 45 km
(Barstow et al. 2011; Tellmann et al. 2009). At this altitude sulphuric acid becomes thermo-
dynamically unstable but fine lower haze of presumably different composition exists down
to ∼30 km. There is also a tentative indication of near-surface haze (Grieger et al. 2004)
whichisnotshowninFig.37.
Cloud Morphology Morphology of the cloud tops was recently studied by extensive
imaging of the planet in the near-UV range by Venus Express and Akatsuki orbiters. The
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Clouds and Hazes of Venus Page 51 of 61 126
Fig. 37 Sketch of the Venus
clouds latitude-altitude structure:
upper haze (light blue), upper
cloud (blue and yellow
“droplets”), middle and lower
clouds (blue and red “droplets”),
lower haze (red “droplets”).
Black contours show isolines of
the Venera-15 temperature field
(from Lellouch et al. 1997). Blue,
green and red solid lines mark the
cloud top (τ=1level),the
tropopause and the cloud base
respectively. Mean parameters of
the cloud layers are shown in
semi-transparent inserts
observations revealed in unprecedented detail and variability the well-known UV markings
created by non-uniform distribution of the unknown absorber or absorbers. The low latitudes
are dominated by mottled and fragmented clouds clearly indicating vivid convective activity
centred in the sub-solar region (Figs. 13–18; Titov et al. 2012). The aerosol scale height in
the upper cloud derived from IR observations is quite large here: 3.5–5 km (Figs. 3,11). The
low and middle latitudes are significantly darker at UV wavelengths than the polar regions
suggesting continuous supply of UV-dark material from depth. The dark equatorial band
whose shape resembles a “Y” or “V” letter is likely to be created by this process.
At 50–60◦latitude the mottled cloud pattern gives way to streaky morphology suggesting
that horizontal, almost laminar, flow prevails here. The dark mid-latitude band has very sharp
poleward boundary likely indicating strong jet-like flows right at the cloud tops. The streaks
inclination with respect to the latitudinal circles is consistent with the measured zonal and
meridional wind speeds.
At high latitudes the planet is covered by almost featureless UV-bright polar hood. Size
of the “polar cap” varies, but usually it encompasses the “cold collar” and polar regions and
approximately coincides with the cloud top depression. The “polar cap” appearance changes
from transparent haze through which the main cloud is clearly visible to dense, bright and
featureless hood (Fig. 17). The “polar cap” is often crossed by thin dark circular or spiral
“grooves” a few hundred kilometres wide that are likely to be created by local jet flows. The
most frequently observed polar feature is the UV-dark oval located at ∼70◦S. The “eye” of
the hemispheric vortex—a strongly variable bright feature observed in the thermal-IR—is
located inside the oval and is dynamically related to it. Since the cloud top is the deepest at
the pole, it has the highest temperature here. After the Pioneer Venus mission this feature
was dubbed “polar dipole”, but recent imaging by Venus Express showed that the shape of
the vortex eye can change from a simple oval to complex multi-centre structure (Fig. 19). Li-
maye et al. (2009) succeeded to simulate the “eye” morphology with an idealized nonlinear
and non-divergent barotropic model and found similar structures in the modelled vorticity
field. In general, the hemispheric vortex circulation on Venus has remarkable morphological
similarities to the Earth’s tropical cyclones and hurricanes, although the energy sources in
two cases are different.
The latitudinal behavior of the cloud top altitude and aerosol structure is correlated with
the mesospheric temperature field (Figs. 8,37). The cloud top starts descending with latitude
at the equatorward edge of the “cold collar” (∼60◦) and approximately follows 220–230 K
temperature isolines (Fig. 37). This trend is accompanied by a decrease of the aerosol scale
height reaching less than 1 km value in the “cold collar”. At higher latitudes the aerosol
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126 Page 52 of 61 D.V. Titov et al.
scale height is usually less than 2 km, but in some warm locations clouds may be rather
diffuse with scale height exceeding 5 km (Fig. 11). This indicates strong variability of the
cloud and temperature structure in the polar regions.
The cloud vertical structure and morphology can be correlated with the observed temper-
ature and dynamic features. Drastic changes in the cloud morphology, i.e. transition from
patchy and mottled clouds in low latitudes to bright uniform polar hood, occurs in the “cold
collar” region (50–60◦). We suggest that this behavior is likely related to the Hadley cir-
culation in the mesosphere that has upwelling branch at equator and downwelling motions
at ∼60◦latitude. Very “sharp” cloud top with small aerosol scale height (≤2 km) in the
“cold collar” coincides with the coldest regions in temperature inversions (Figs. 8,37). This
feature is likely of radiative origin. The sharp cloud top boundary provides effective cooling
to space that maintains low temperatures. Strong temperature inversions in turn create con-
vectively stable conditions above the cloud top thus preventing vertical mixing of aerosols
and maintaining sharp cloud boundary. This negative feedback between dynamical condi-
tions and radiative effects maintains a stable cloud top structure in the “cold collar” region.
The middle and high latitudes are also characterized by vanishing vertical wind shear within
the clouds (Sánchez-Lavega et al. 2008) that suppresses wind shear instabilities. Titov et al.
(2008) used the peculiarities of stability distribution at the cloud top to qualitatively explain
the global pattern of the UV markings.
High resolution imaging of the Venus cloud top enabled detailed study of small scale
features (Fig. 18). In low latitudes the mottled clouds form convective cells and wave trains
with typical scale of 100–200 km. “Column” clouds resembling Earth cumulus clouds few
tens of kilometres across are often present. These structures as well as small streaks show
signs of strong lateral advection right above the cloud top.
Imaging in the near-IR transparency “windows” on the night side showed significant vari-
ations of total cloud opacity thus revealing morphology of the middle and lower cloud. The
general trend is that the cloud is the most transparent at 40–60◦latitude with opacity increas-
ing towards the equator and the poles (Crisp et al. 1991). Poleward of ∼60◦the cloud be-
comes so opaque that almost no radiation can escape to space and no information on the deep
cloud morphology can be derived from the observations (Fig. 20). The models that couple
general circulation and cloud formation processes imply that meridional Hadley circulation
can lead to opacity increase in both ascending branch at equator and descending flow at the
poles (Imamura and Hashimoto 1998). The deep cloud has patchy or blocky structure equa-
torward of ∼40◦. Opacity inhomogeneities can be as large as 2000 km (Fig. 12). Lifetime of
the holes spans from few hours to several weeks with larger inhomogeneities living longer.
At middle latitudes the features become elongated and tilted with respect to latitudinal cir-
cles that is consistent with vortex circulation converging at the poles (Belton et al. 1991;
McGouldrick et al. 2008). The overall pattern of opacity inhomogeneities observed in the
deep cloud suggests active convection equatorward of ∼40◦and laminar circulation at
higher latitudes similar to what is seen at the cloud top.
Composition and Chemistry The Venus cloud is a complex chemical system in tight
interaction with the gaseous chemistry and photochemistry of the atmosphere. Both obser-
vations and models provide strong evidence that at least aerosols of the upper cloud and
upper haze are composed mainly of concentrated sulphuric acid H2SO4with mass concen-
tration of 75–83% at the cloud tops (Esposito et al. 1997; Cottini et al. 2015; Krasnopolsky
2015) with tentative trend to increase up to 90% in the polar regions (Barstow et al. 2011).
Sulphuric acid is formed from sulphur dioxide and water vapour under solar light. The pho-
tochemical “factory” is located in the upper cloud layer (Figs. 27,37). Nature of additional
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Clouds and Hazes of Venus Page 53 of 61 126
components responsible for the near-UV absorption in the upper cloud is still a matter of
debate but the most favourable candidates are iron chloride and sulphur allotropes or related
sulphur species. Observations also suggest presence of more than one UV-blue absorbers.
Sulphuric acid evaporates at the cloud bottom near 48 km and is thermally decomposed at
about 40 km (Krasnopolsky 2013).
Several experiments on descent probes indicated presence of non-sulphuric acid aerosols
in the deep atmosphere. In-situ chemical analysers onboard Venera and Vega probes de-
tected presence of chlorine, phosphorous and iron compounds in relatively large amounts
in addition to sulphur bearing particles. The optical measurements onboard Pioneer Venus
tentatively suggested the presence of large (r>3μm) non-spherical particles for which
sulphuric acid composition is excluded. The lower hazes detected at 33–48 km by several
descent probes are found at temperatures too elevated for sulphuric acid droplets to ex-
ist.
Microphysical Properties Venus clouds consist of fine particles with sizes of up to about
10 μm. According to LCPS/Pioneer Venus cloud particle spectrometer (Knollenberg and
Hunten 1980) the particle size distribution in the main cloud deck has up to three modes
(Esposito et al. 1983). Mode 1 includes fine sub-micron aerosols (r∼0.1–0.2μm) with
number density of few thousand particles per cm3(Fig. 37). Mode 2 consists of micron size
particles (r∼1μm) with number density n∼50 cm−3. Modes 1 and 2 are ubiquitously
present in the main cloud deck. Larger particles (mode 3) with r∼3–4 μm and number
density of few tens of particles per cm−3were found in the middle and lower cloud layers
below 56 km. However, the controversy whether large particles form a separate mode or they
are just a “tail” of the mode 2 is still not resolved (Esposito et al. 1997). Also ISAV-A particle
size spectrometer onboard Vega-1, -2 probes that descended close to equator detected very
small number of particles with r>2.5μm and did not confirm existence of mode 3 aerosols
(Moshkin et al. 1986; Gnedykh et al. 1987).
All observations suggest that the cloud properties in the “cold collar” and polar regions
significantly differ from those at low latitudes. Analysis of the near-infrared emissions on the
Venus nightside revealed anomalous particles in the deep cloud in the polar regions (Wilson
et al. 2008; Barstow et al. 2011;Hausetal.2013). These particles are either larger or differ
in composition from those elsewhere on the planet.
The upper haze is ubiquitously present in the mesosphere at altitude 70–100 km. Micro-
physical properties of the aerosols in this region were extensively investigated in solar and
stellar occultations onboard Venus Express (Wilquet et al. 2009; Luginin et al. 2016). Both
submicron and micron size particles were found here. Gao et al. (2014) reproduced the ob-
served size distribution and number density of aerosol particles in the main cloud and upper
haze using coupled microphysical and vertical transport models. Interestingly, the model
generated a quasi-periodically varying system rather than a system with a stable equilibrium
distribution. The upper haze properties appear to result from a mixture of droplets formed
from in situ nucleation of sulphuric acid vapor on meteoric dust and droplets upwelled from
the cloud decks below.
10 Open Issues and Future Studies
The past couple of decades have seen remarkable progress in our understanding of the Venus
cloud system. This was achieved mainly by remote sensing techniques that allowed detailed
characterization of the morphology and vertical structure of the cloud layer including its
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126 Page 54 of 61 D.V. Titov et al.
spatial and temporal variations. The most important remaining open issues are related to the
composition and chemistry of the cloud layer. The detection of chlorine and phosphorous
compounds in amounts exceeding that of sulphur by the Vega descent probes still challenges
chemical models. Moreover, the nature and origin of the unknown UV-blue absorbers that
are of crucial importance for understanding of chemistry, radiative balance and dynamics of
the Venus atmosphere still remains a mystery. This justifies the urgent need for an in-situ
analysis of Venus aerosols onboard descent probes and especially long-living aerial plat-
forms (balloons) with a capability of changing floating altitude in order to unambiguously re-
solve both open issues in composition of the Venus clouds. Middle-to-high-resolution Venus
spectra would also be welcome, particularly if they provide spatial resolution to separate re-
gions with higher and lower UV-blue absorption. The wavelength range of 0.4–0.5μmisof
the highest interest, as it may provide very useful constraints on the nature of the near-UV
absorption in Venus atmosphere. High-resolution spectra would also help to separate spec-
tral features of gaseous and particulate species, thus constraining the physical state of the
near-UV absorbers. High-resolution spectroscopy in the near-IR transparency “windows”
still have a great potential in characterization of the atmospheric chemistry below the clouds
that is strongly related to the clouds composition.
The Venus clouds appear to be a very variable system with many complex couplings
that are far from being understood. Spatial and temporal variations of the planet albedo,
especially in the UV-blue domain, as well as total cloud opacity need to be investigated to
detect changes over various time scales. These studies are of great importance for under-
standing of energy balance and dynamics of the Venus atmosphere. Measurements in the
cloud layer are needed to distinguish between inorganic and hypothetic organic/biologic
absorbers.
These efforts should be supported by laboratory modelling of aerosol chemistry and study
of optical properties of its particulate products. As future work, there are essentially two as-
pects that should be investigated. Laboratory spectra of already proposed or new candidates
for the near-UV absorption at the conditions of the upper Venus atmosphere (temperature,
pressure, and solar radiation) are required. Recent years witnessed development of sophis-
ticated microphysical models simulating the processes of formation, growth and decompo-
sition of sulphuric acid aerosols. These models should be expanded to the other species (for
instance, elemental sulphur) that are the candidates for the unknown near-UV absorbers. The
next step would be incorporation of the microphysical models in regional and global circu-
lation models that would elucidate the feedbacks between microphysics, chemistry, and the
momentum and energy balance.
Laboratory studies of sulphuric acid aerosols at high concentrations including phase be-
haviour are required in order to better constrain the microphysical properties of the aerosols
under Venus conditions. The relative solubility between the sulphuric acid and potential
cloud condensation nuclei has tremendous effects on the size and mass distributions of the
Venus aerosol. The laboratory support is of high importance for developing consistent mod-
els of the cloud evolution.
In order to improve our understanding of the plausibility and, if it exists, the characteris-
tics of Venus lightning, additional observations, as well as detailed laboratory and modelling
studies are required. Search for optical detection of lightning at the cloud tops by orbiting
spacecraft (Borucki et al. 1996; Takahashi et al. 2008) and in situ optical, radio wave and
electric field monitoring can help elucidate the frequency and location of the occurrence of
Venus lightning. These efforts should be complemented by laboratory studies exploring the
charging properties of multi-phase sulphuric acid in Venus-like conditions, or of homoge-
neous triboelectric charging of glassy sulphuric acid are required.
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Clouds and Hazes of Venus Page 55 of 61 126
Acknowledgements This work significantly benefited from discussions within the Venus Clouds Interna-
tional Team supported by the International Space Science Institute in Bern, Switzerland. The authors are
grateful to both reviewers for careful reading of the manuscript and useful comments and suggestions.
Open Access This article is distributed under the terms of the Creative Commons Attribution 4.0 Inter-
national License (http://creativecommons.org/licenses/by/4.0/), which permits unrestricted use, distribution,
and reproduction in any medium, provided you give appropriate credit to the original author(s) and the source,
provide a link to the Creative Commons license, and indicate if changes were made.
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