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date: 19 August 2018

The Surface of Venus

Summary and Keywords

This chapter reviews the conditions under which the basic landforms of Venus formed, interprets their nature, and analyzes their local, regional, and global age relationships. The strong greenhouse effect on Venus causes hyper-dry, almost stagnant near-surface environments. These conditions preclude water-driven, and suppress wind-related, geological processes; thus, the common Earth-like water-generated geological record of sedimentary materials does not currently form on Venus. Three geological processes are important on the planet: volcanism, tectonics, and impact cratering. The small number of impact craters on Venus (~1,000) indicates that their contribution to resurfacing is minor. Volcanism and tectonics are the principal geological processes operating on Venus during its observable geologic history.

Landforms of the volcanic and tectonic nature have specific morphologies, which indicate different modes of formation, and their relationships permit one to establish their relative ages. Analysis of these relationships at the global scale reveals that three distinct regimes of resurfacing comprise the observable geologic history of Venus: (1) the global tectonic regime, (2) the global volcanic regime, and (3) the network rifting-volcanism regime. During the earlier global tectonic regime, tectonic resurfacing dominated. Tectonic deformation at this time caused formation of strongly tectonized terrains such as tessera, and deformational belts. Exposures of these units comprise ~20% of the surface of Venus. The apparent beginning of the global tectonic regime is related to the formation of tessera, which is among the oldest units on Venus. The age relationships among the tessera structures indicate that this terrain is the result of crustal shortening. During the global volcanic regime, volcanism overwhelmed tectonic activity and caused formation of vast volcanic plains that compose ~60% of the surface of Venus. The plains show a clear stratigraphic sequence from older shield plains to younger regional plains. The distinctly different morphologies of the plains indicate different volcanic formation styles ranging from eruption through broadly distributed local sources of shield plains to the volcanic flooding of regional plains. The density of impact craters on units of the tectonic and volcanic regimes suggests that these regimes characterized about the first one-third of the visible geologic history of Venus. During this time, ~80%–85% of the surface of the planet was renovated. The network rifting-volcanism regime characterized the last two-thirds of the visible geologic history of Venus. The major components of the regime include broadly synchronous lobate plains and rift zones. Although the network rifting-volcanism regime characterized ~2/3 of the visible geologic history of Venus, only 15%–20% of the surface was resurfaced during this time. This means that the level of endogenous activity during this time has dropped by about an order of magnitude compared with the earlier regimes.

Keywords: Venus, tectonics, volcanism, stratigraphy, regimes of resurfacing

Venus is one of the terrestrial planets (those that are inside of the orbit of Jupiter) and by its general planetary parameters is very similar to Earth (Table 1a&1b). The mean radius of Venus is about 6,051.8 km (~0.95 of Earth), its mass is about 4.87 × 1024 kg (~0.8 of Earth), and the mean density is about 5.24 g/cm3 (~0.95 of Earth). Because of these similarities, for a long time, Venus and Earth were considered as twins. Both planets are much larger than the other terrestrial planets such as Mercury, Moon, and Mars (Figure 1).

Table 1a. Parameters of the Terrestrial Planets

Planet

Mass, 1024 kg

Radius, km

Volume, 109 km3

Density, g/cm3

Distance from Sun, 106 km

Mercury

0.33

2,440

61

5.43

57.91

Venus

4.87

6,052

928

5.24

108.2

Earth

5.97

6,371

1,083

5.51

149.6

Moon

0.07

1,737

22

3.35

149.6

Mars

0.64

3,390

163

3.92

227.9

Table 1b. Parameters of the Terrestrial Planets Relative to Earth

Planet

Mass

Radius

Volume

Density

Distance from Sun, AU

Mercury

0.06

0.38

0.06

0.98

0.39

Venus

0.82

0.95

0.86

0.95

0.72

Earth

1

1

1

1

1

Moon

0.01

0.27

0.02

0.61

1

Mars

0.11

0.53

0.15

0.71

1.52

(*) Astronomical Units, 149.6 km.

The Surface of VenusClick to view larger

Figure 1. Comparative sizes of the terrestrial planets illustrated by their volumes.

The relatively small dimensions of these planets have caused termination of their evolution at earlier stages. The testimony of this is the rich impact crater record on these planets formed by an external process of meteorite bombardment. Its intensity was much greater near the beginning of planetary history, but rapidly declined. The large number of impact craters on the surfaces of Mercury, Moon, and Mars indicate that internally generated planetary processes and activity, such as volcanism and tectonics, were subordinate to impact cratering. The rich crater records indicate that the volcanic and tectonic activity on the smaller terrestrial planets faded away earlier in their planetary history and has not continued at a high level after the waning of the meteorite bombardment. In contrast, the crater records on Venus and Earth are poor, which means that internal processes on these planets continued for a much longer time; they are active on Earth today, and perhaps on Venus too. Thus, these planets represent examples of the later planetary evolution.

Despite the large-scale similarity to Earth, the surface conditions on Venus differ greatly from those on our planet. The atmosphere of Venus consists mostly of carbon dioxide (~97%) and the atmospheric surface pressure is nearly 100 times that of Earth (or equivalent to pressures at about 1 km depth in the sea). The temperature at the surface is high, about 770°K. Rotation of Venus is slow and retrograde, and a single Venusian day is 243 Earth days. Unlike on Earth, there is no evidence for a magnetic field on Venus, which is either a consequence of slow rotation of the planet or a reflection of fundamental differences in the core composition or size.

Thus, Venus and Earth are not twin planets but, nonetheless, represent advanced stages of the planetary evolution. Why are planets that are so similar in their major parameters (Table 1a&1b) so environmentally different? What major factors have caused their diversity? Does Venus represent a possible fate of Earth? Do both planets follow their principally different paths of evolution? These problems, at the current level of our knowledge of the exoplanets, can be solved only by the mean of comparative studies of Venus and Earth.

Among the terrestrial planets, Venus is the most difficult object to investigate. The dense, opaque atmosphere, permanently capped with clouds, requires radar-based techniques to study the geological structure on the surface of Venus. During the earlier Earth-based radar observations of Venus from the Goldstone Observatory, several radar-bright and elevated features, named Alpha and Beta, were detected. The spatial resolution of these images was rather low, about 10 km. They showed the presence on the surface of various circular and elongated bright and dark features, the nature of which remained unclear.

The radar system of the Arecibo Observatory has a greater sensitivity and allows mapping of Venus at higher resolution (5–20 km) within a larger area comprising ~25% of the surface. Images taken at the Arecibo Observatory have revealed a number of brighter and darker features, among which only a few are circular and resemble impact craters. The scarcity of the crater-like forms led to the conclusion that the surface of Venus is relatively young in comparison with the surfaces of the smaller terrestrial planets (Campbell & Burns, 1980). It was a discovery of fundamental importance, because it indicated that the balance between the external (impact cratering) and internal (volcanism, tectonics) processes of resurfacing on Venus is shifted toward endogenous activities.

The major limitation of Earth-based observations of Venus is that they are restricted to about the same region of the planet, between ~60°S and 75°N and from ~260°E through zero meridian to ~30°E. Several interplanetary missions to Venus have greatly improved this situation. During these missions (Table 2), the surface of the planet has been studied both from orbit and from the surface.

Table 2. Successful Missions to Venus

Mission

Arrival

Termination

Objective

NASA USA, Mariner 2

December 1962

January 1963

Flyby

USSR, Venera 4

October 1967

October 1967

Lander

NASA USA, Mariner 5

October 1967

Nov 1967

Flyby

USSR, Venera 5

May 1969

May 1969

Atmospheric probe

USSR, Venera 6

May 1969

May 1969

Atmospheric probe

USSR, Venera 7

December 1970

December 1970

Lander

USSR, Venera 8

July 1972

July 1972

Lander

NASA USA, Mariner 10

February 1974

March 1975

Flyby

USSR, Venera 9

October 1975

October 1975

Lander

USSR, Venera 10

October 1975

October 1975

Lander

NASA USA, Pioneer Venus 1

December 1978

August 1992

Orbiter

NASA USA, Pioneer Venus 2

December 1978

December 1978

Bus

Large probe

North probe

Night probe

Day probe

USSR, Venera 13

March 1982

March 1982

Lander

USSR, Venera 14

March 1982

March 1982

Lander

USSR, Venera 15/16

October 1983

July 1984

Orbiter

USSR, Vega 1

June 1985

January 1987

Flyby

June 1985

June 1985

Lander

June 1985

June 1985

Balloon

USSR, Vega 2

June 1985

March 1987

Flyby

Lander

Balloon

NASA USA, Magellan

August 1990

October 1994

Orbiter

ESA, Venus Express

Apr 2006

December 2014

Orbiter

Lander Studies of Venus

The primary objective of the landers was the collection of geochemical information. Our knowledge of the geochemistry of solid Venus was, and remains, very limited. This is mostly because the dense layer of clouds prevents spectral studies of the planet. Thus, landers are the only means by which geochemical data from the surface of Venus can be obtained. Several landers (Table 2) visited the planet in the period from 1972 (Venera-8) to 1985 (VEGA-1 and 2) and reported the first and only data on the chemical composition of soils on the surface of Venus. Selection of the landing sites (Figure 2) was based purely on interplanetary ballistic constraints, because little was known about the surface geology when the Venera-VEGA missions were planned and implemented.

The Surface of VenusClick to view larger

Figure 2. Map showing landing sites of the Soviet landers of the Venera and Vega missions. Background is a mosaic (from 82.5°N to 82.5°S) of the SAR (synthetic aperture radar) images taken during the Magellan mission. The map is in simple cylindrical projection.

In four landing points (Venera-8, -9, -10, and VEGA-1), concentrations of the three major thermal-generating components, K, Th, and U, were determined by gamma spectrometry (Table 3) (Surkov, 1997). The mean values of their concentrations on Venus are well within the range that is typical of terrestrial basalts. However, the enhanced concentrations of K, Th, and U in soils at the Venera-8 landing site made it possible to interpret the results of this station as evidence for the presence of a non-basaltic material on Venus.

Table 3. Summary Data on Chemical Composition of Materials on the Surface of Venus

Lander

Unit

SiO2

TiO2

Al2O3

FeO

MnO

MgO

CaO

K2O

S

Cl

K

Th

U

V-8

psh

4.0±1.2

6.5±0.2

2.2±0.7

V-9

rp1

0.5±0.1

3.7±0.4

0.6±0.2

V-10

rp1

0.3±0.2

0.7±0.3

0.5±0.3

V-13

rp1

45.1±3.0

1.59±0.45

15.8±3.0

9.3±2.2

0.2±0.1

11.4±6.2

7.1±0.96

4.0±0.63

0.65±0.4

<0.3

V-14

pl

48.7±3.6

1.25±0.41

17.9±2.6

8.8±1.8

0.16±0.08

8.1±3.3

10.3±1.2

0.2±0.07

0.35±0.31

<0.4

Vg-1

rp1

0.45±0.22

1.5±1.2

0.64±0.47

Vg-2

rp2

45.6±3.2

0.2±0.1

16.0±1.8

7.74±1.1

0.14±0.12

11.5±3.7

7.5±0.7

0.1±0.08

1.9±0.6

<0.3

0.40±0.20

2.0±1.0

0.68±0.38

Note: All major elements are in wt. %, Th and U are in ppm; V—landers of the Venera series, Vg—landers of the Vega series

In two landing sites (Venera-13, and -14), the concentrations of the major petrogenic oxides (without Na2O) were measured by the X-ray fluorescence (XRF) method (Table 3) (Surkov, 1997). In one point (VEGA-2), both methods (gamma spectrometry and XRF) were used separately, and the concentrations of the thermal-generating elements and the major oxides were measured (Table 3). The XRF data also suggested that rocks of basaltic composition make up the surface in the landing sites.

Two important factors strongly limit the value of the Venera/VEGA data and prevent their robust interpretation. The first major problem is that we do not know the exact position of the landers and it is impossible to understand which type of materials the landers investigated. All stations landed somewhere within their own landing circle, which is ~300 km in diameter and embraces terrains of different origin and age. For example, the landing circle of Venera-10 includes six various and extensive units related to tectonics and volcanism. According to the Venera-10 panorama (Figure 3), the station is on a flat, sub-horizontal surface. This type of surface favors vast volcanic plains as the host units for the lander and disfavors tectonized units, although they cannot be ruled out confidently.

The Surface of VenusClick to view larger

Figure 3. A part of a panorama taken by the Venera-10 lander. Bright features near the lower edge of the image are construction details of the lander (up to 10–15 cm long). The image shows a surface with plates of bare rock (light tones) overlaid by accumulations of loose soil (darker tones).

Thus, at the landing points, association of the chemical data with the specific terrains can be made on a probabilistic basis only (Abdrakhimov, 2005). The lack of knowledge of the exact location of the landers prevents understanding of the geochemical aspects of nature of the units and, thus, prevents formulation of reasonable petrogenic models.

The second and most important limitation of the Venera/VEGA geochemical data is the low precision of the measurements made on the surface (Table 3). The large error bars keep the data collected by the landers at a low level of possible interpretation. For example, on a ternary plot that shows relationships of the major thermal-generating elements (Figure 4), points of terrestrial magmatic rocks form a prominent trend that reflects broad variations of Th/K ratio. The points of the Venera-8 and -9 landers, which have the smallest errors, fall onto the terrestrial trend. The mean values of Venera-10 and VEGA-1 and -2 seem to be shifted from the main terrestrial trend toward the U-side of the diagram (Figure 4). However, the error ellipses of these measurements are so large that such interpretations are not well constrained.

The Surface of VenusClick to view larger

Figure 4. A ternary diagram that shows relative concentrations of the major thermal generating components (potassium, thorium, and uranium) for the terrestrial volcanic rocks from different geodynamic environments. Stars indicate measurements made on the surface of Venus by landers of the Venera (V) and Vega (Vg) missions. Black and white ellipses around the stars indicate the measurement error bars.

Orbiter Studies of Venus

The primary goals of the orbiters were mapping of the surface at the global scale and collecting of data on morphology, topography, and gravity of Venus.

Pioneer Venus (PV) was the first orbiter that systematically observed the entire surface of Venus (Table 2). These data allowed compilation of a global radar map of the planet, the spatial resolution of which was still low, about 30 km. The low resolution prevented a confident geological interpretation of features detected on the surface. The topographic data collected by the PV orbiter have a vertical accuracy of ~200 m and very coarse spatial resolution, ~75–100 km. Despite their low resolution, the topographic data have indicated that the hypsogram of Venus is unimodal, which is in a sharp contrast with the bimodal hypsogram of Earth (Figure 5). These differences in the global topography distribution reflect fundamental differences between Venus and Earth and are related to both the crustal structure and the dominant tectonic styles on both planets.

The Surface of VenusClick to view larger

Figure 5. Hypsograms of Earth (dashed line) and Venus (solid thick line) that show a fraction of the planetary surface within a specific elevation interval. The curves show a drastic difference in the distribution of global topography on both planets. The bimodal hypsogram of Earth reflects the different topographic configurations of oceanic (left peak) and continental (right peak) lithosphere. The hypsogram of Venus is strongly unimodal.

Plate tectonics and continental accretion on Earth constantly form and separate in space two principal components of Earth’s crust, continental and oceanic, and keep the oceanic lithosphere thin and hence low-lying relative to the continents. If plate tectonics were to be operating on Venus, the unimodal hypsogram would be more consistent with the crust that is made of one component of about the same density. If plate tectonics do not work on Venus, its crust could be single- or poly-component, but the unimodal hypsogram would indicate that the components are overlapping each other, as occurs on Moon.

The next radar-bearing orbiters at Venus were those of the Venera-15/16 (V-15/16) mission (Table 2). During this mission, about the upper half of the northern hemisphere of Venus was mapped at unprecedented high spatial resolution, 1–2 km (Barsukov et al., 1986). The incidence angle of the radar systems (the angle between the vertical and radar beam) to a great degree relates the strength of the returned signal and the scale of the surface roughness. The incidence angle of the V-15/16 radars was ~10°, and the returned signal was preferentially modulated by the small-scale (a few kilometers) topographic variations. The images taken by V-15/16 radars allowed both the inventory and the genetic interpretation of features seen on the surface. After the V-15/16 mission, it become clear that smooth (presumably volcanic) plains and rough, tectonized terrains make up the absolute majority of the surface of Venus. Impact craters were unambiguous features in the V-15/16 images, and their small number has confirmed the conclusion about the geological youth of the Venusian surface.

The low spatial resolution of the PV images and incomplete survey of the surface by the V-15/16 mission left the majority of the Venus surface a geological “terra incognita.” This ambiguity was eliminated after the Magellan mission (Saunders et al., 1992). The Magellan radar system mapped virtually the entire surface of Venus (~97%) at a resolution of ~75–100 m. Images taken by the Magellan radar system (SAR, synthetic aperture radar) showed no evidence for the structures that characterize plate tectonics on Earth (Solomon et al., 1992) but a very wide spectrum of volcanic landforms on the surface of Venus (Head et al., 1992). The incidence angle of the Magellan radar beam was larger than that of the V-15/16 mission, and the strength of the reflected signal was modulated mostly by the small-scale (meter-decameter) roughness of the surface. The geometry of the Magellan observations enhanced the visibility of morphologically homogenous terrains and allowed their confidential delineation and the determination of their relative ages. The systematic geological mapping of the Magellan images resulted in the compilation of a global geological map of Venus (Ivanov & Head, 2011) that shows the distribution of the mapped terrains (units) in space and time (Figure 6a, b).

The Surface of VenusClick to view larger

Figure 6. The surface geology of Venus. (a) The geological map shows the spatial distribution of the mapped morphological units. The map is in equal-area Mollweide projection. (b) The correlation chart shows their distribution in time.

Main Type of Terrains Composing the Surface of Venus

The massive CO2 atmosphere of Venus has caused a strong greenhouse effect, which limits the range of possible geological processes operating on the planet. The high surface temperature and apparently low temperature/pressure gradients in the lower atmosphere cause the hyper-dry, almost stagnant near-surface environment, and the common Earth-like water-related geological record of sedimentary materials cannot readily form on Venus. Three geological processes are important on the planet: volcanism, tectonism, and impact cratering. Only about a thousand impact craters have been detected on Venus. This means that: 1) the surface of the planet is relatively young (the mean age estimates vary from ~750 to ~300 Ma), and 2) the contribution of impact craters to resurfacing is minor. Thus, volcanism and tectonics have been the principal geological processes during the observable geologic history of Venus.

Volcanic Landforms

The dimensions of the occurrences, topographic configuration, and relative importance of features of volcanic and tectonic origin permit the subdivision of the volcanic landforms of Venus into three broad categories: (1) regionally to globally distributed volcanic plains, (2) local to regional volcanic features, and (3) local to regional features with both volcanic and tectonic components (volcano-tectonic features).

Volcanic Plains

This category includes vast plains-like surfaces that are either mildly deformed by tectonic structures or nondeformed (Figure 7).

The Surface of VenusClick to view larger

Figure 7. Examples of major volcanic plains on Venus. (a) Numerous small shield- and cone-like mounds interpreted as volcanic constructs populate the surface of shield plains (psh); center of the image is at 38.7°N, 117.3°E. (b) Two subunits of regional plains (rp1 and rp2) typically have different radar albedo that helps to distinguish these types of plains. The brightness of the lower subunit (rp1) is moderate and uniform, and the surface of the upper subunit (rp2) is brighter. A narrow channel cuts the surface of unit rp1 and appears as a prominent topographic feature (arrow). Material of unit rp2 almost completely fills the channels (arrow); center of the image is at 48.1°N, 161.3°E. (c) Numerous brighter and darker lava flows constitute lobate plains (pl); center of the image is at 10.7°N, 263.9°E. (d) The surface of smooth plains (ps) appears homogeneous and has low to moderate radar brightness; center of the image is at 3.5°N, 308.1°E.

Shield plains (psh, Accruva Formation, Figure 7a): Shield plains consist of two components. The most abundant are numerous small (from a few kilometers up to 10 km across) shield- and cone-like mounds. These features, which often occur in clusters, are likely to represent volcanic edifices. The second component is represented by morphologically smooth plains that host the mounds and have the same radar albedo as the adjacent shields and cones. Shield plains occur as outliers several tens to hundreds of kilometers across (Figure 6a) that tend to be at slightly higher elevations than the surrounding regional plains. The total area of the plains is about 84.5 × 106 km2, or 18.5% of the surface of Venus (Table 4). Analysis of the geology of the Venera-8 landing site has shown that the lander is likely to have sampled the surface of shield plains.

Table 4. Area of Units Mapped on the Global Geological Map

Unit

Area, 106 km2

Percent of the surface

t

33.2

7.3

pdl

7.2

1.6

pr

9.6

2.1

mt

1.3

0.3

gb

37.1

8.1

psh

79.3

17.4

rp1

141.8

31.1

rp2

42.0

9.2

sc

3.3

0.7

rz

22.6

5.0

ps

10.3

2.3

pl

37.8

8.3

craters

2.6

0.5

data gaps

28.1

6.1

TOTAL

456.3

100.0

Note: Total area is between 82.5°N and 82.5°S.

There are three types of settings on Venus where the occurrences of the plains are less abundant or absent (Figure 6a). The first corresponds to regions of thickened crust such as large tessera massifs. Thicker crust in these areas may have served as a rheological or density barrier inhibiting formation of shield plains. The second is in regions of chasmata, in which young and abundant extensional structures may have destroyed exposures of shield plains. The third type occurs within wide lowland areas, in which younger volcanic materials may have buried occurrences of shield plains.

Regional Plains (rp): Regional plains represent the most widespread type of volcanic plains on Venus, covering ~195.5 × 106 km2, or 42.8% of the surface of Venus (Table 4). Regional plains are composed of morphologically smooth and homogeneous plains materials with intermediate-dark to intermediate-bright radar albedo. Narrow wrinkle ridges cut both shield and regional plains and form pervasive intersecting networks. The Venera 9, 10, and 13 and Vega 1 and 2 landers probably landed on the surface of regional plains. Characteristic variations of radar albedo permit the subdivision of regional plains into two subunits.

The lower subunit of regional plains (rp1, Rusalka Formation, Figure 7b) has a surface with relatively low radar albedo. This unit is the most abundant unit on Venus (about 150.7 × 106 km2, or 33.0% of Venus, Table 4). Fields of the lower subunit of regional plains extend for hundreds to thousands of kilometers, connect remote regions, and can be traced almost continuously around the globe (Figure 6a). The unit preferentially overlays the floor of the lowlands, in which occurrences of shield plains are rare; it surrounds the major tessera-bearing uplands and occurs between elevated regions composed of the heavily tectonized units and shield plains (Figure 6a). Sources of the plains are commonly not seen at the resolution of Magellan images.

The upper subunit of regional plains (rp2, Ituana Formation, Figure 7b) has a surface with a higher radar albedo than that of the lower subunit (Figure 7b). The upper subunit of regional plains overlays about 44.8 × 106 km2, or 9.8% of Venus (Table 4), and occurs usually as equidimensional or slightly elongated patches of flow-like shape that are tens to several hundreds of kilometers across (Figure 6a). Fields of the upper subunit of regional plains do not occur within the large tessera regions and tend to avoid large lowland regions, the surface of which is made up by the lower subunit of regional plains (Figure 6a). Occurrences of rp2 surround some of the large volcanic centers such as coronae and large volcanoes and form distal aprons of volcanic materials around them.

Lobate plains (pl, Bell Formation, Figure 7c): Occurrences of lobate plains represent complexes of numerous radar bright and dark flow-like features that superpose each other. Individual flows within the plains can be several hundred kilometers long and tens of kilometers wide. Lobate plains make up a significant portion of the surface of Venus, about 40.3 × 106 km2, or 8.8% of Venus (Table 4), and their occurrences form equidimensional fields from many tens and up to many hundreds of kilometers across (Figure 6a). Lobate plains characterize many large volcanic centers on Venus. The plains are usually associated with the large dome-shaped topographic rises, whereas the elevated plateau-shaped tessera regions lack significant occurrences of lobate plains. Within the Beta-Atla-Themis (BAT) region, lobate plains are spatially associated with rift zones. The Venera-14 lander probably landed on the surface of one of the occurrences of lobate plains.

Smooth plains (ps, Gunda Formation, Figure 7d): These plains have a smooth, tectonically undisturbed, and featureless surface. Areas of smooth plains are usually characterized by a low radar albedo and appear as dark spots on the surface. This unit makes up a small portion of the surface, about 10.3 × 106 km2, or 2.3% of Venus (Table 4). There are three types of geological settings of smooth plains: (1) near and within volcanic regions, where the plains are closely associated with lobate plains; (2) dark plains in spatial association with impact craters that are likely to represent remnants of dark parabolas formed due to impact events; (3) patches of smooth plains within some large tessera regions that probably have a volcanic origin.

Volcanic Features

This category includes shield volcanoes, steep-sided domes and festoons, and narrow sinuous channels (Figure 8).

The Surface of VenusClick to view larger

Figure 8. Examples of volcanic features on Venus. (a) A large shield volcano on Venus. The summit of the volcano (center of the image) shows the presence of a relatively small caldera-like feature, from which numerous flows of lobate plains extend radially. The diameter of the volcano is about 400 km and its height about 2 km. Thus, the average slope on the flanks of the volcano is about 0.6°; center of the image is at 9.3°N, 29.3°E. (b) Steep-sided domes (arrows). The group of domes is spatially associated with a corona and a branch of a groove belt extending northward of the corona. Edges of some of the domes are collapsed/fluted and embayed by regional plains; center of the image is at 32.1°N, 312.3°E. (c) The southwestern portion of a large complex of viscous flows (festoon). The festoon was built by multiple eruptions of lavas that formed broad fan-shaped lobes with a steep-sided frontal scarp; center of the image is at 38.1°S, 163.9°E. (d) A narrow, long, and sinuous channel cuts the surface of the lower subunit of regional plains (rp1); center of the image is at 44.9°N, 15.8°E.

The shield volcanoes (Figure 8a) represent significant (up to a few kilometers) topographic highs with a radiating pattern of flows. The flows belong to either the upper subunit of regional plains (73 volcanoes, or ~33% of their population), lobate plains (115 volcanoes, ~53%), or both (28 volcanoes, ~13%). No large volcanoes were detected in association with either shield plains or the lower subunit of regional plains. Diameters of the shield volcanoes vary over a wide range, from ~20 to ~1,000 km. Those volcanoes that occur as a part of the upper subunit of regional plains are smaller, about 50–250 km across, in comparison with the volcanoes associated with lobate plains, which are about 300–600 km in diameter.

Steep-sided domes and festoons (Figure 8b, c) Step-sided domes (Figure 8b) are rounded features with a flat top surface and a prominent frontal scarp. The domes are usually many hundreds of meters high and tens of kilometers in diameter; the most common diameters are 10–30 km. About 320 steep-sided domes were mapped on Venus; most of them are spatially associated with occurrences of shield plains and they tend to avoid the large expanses of regional plains. Detailed analysis of the stratigraphic position of the domes has shown that the materials of regional plains embay domes. The frontal scarp of the domes suggests that their materials upon eruption were more viscous than the usual basaltic lava, which is consistent with either more silicic composition of the magma, its higher degree of crystallization, or abundant gas bubbles in the erupting material.

Eruption of the highly crystallized magma requires a delicate balance between the time of thermal evolution of magma in a chamber, injection of new batches of magma, and eruption of its crystallized portion. Although emplacement of basaltic lavas with high concentration of crystals is known on Earth, it represents a unique phenomenon and can hardly explain the numerous steep-sided domes on Venus that have a wide range of diameters.

Eruption of lavas with abundant gas bubbles is a more likely explanation, because the high atmospheric pressure on Venus would tend to lock volatiles in the rising magma and cause eruption of a viscous lava “foam.” If the high ambient pressure were the primary factor of the domes formation, they would tend to occur at the lower elevations, where the atmospheric pressure is higher. In contradiction to this prediction, the topographic distribution of the domes is shifted to the right, and they tend to occur at the middle topographic levels. Thus, higher SiO2 content in the magma appears to be the most plausible explanation of the mode of formation of the steep-sided domes, and they are considered the main morphologic evidence of non-basaltic volcanism on Venus.

There are three large (up to a few hundred kilometers across) volcanic flows on Venus that are characterized by pronounced frontal scarps (Figure 8c) and are likely to have been formed by eruptions of lavas that apparently were more viscous than basalt. In contrast to the domes, however, festoons have a highly irregular planimetric shape. Two of the festoons are within the lower-lying territories; the third is on top of high-standing Ovda Regio. This topographic distribution of festoons provides additional evidence that the ambient atmospheric pressure did not play a major role in the changing of the effective viscosity of materials of the steep-sided domes.

Long, narrow, and sinuous channels (Figure 8d) represent another type of feature that is likely to have a volcanic origin. The typical width of the channels is about 0.5–1.5 km, and their length can be many hundreds of kilometers. The longest channel on Venus, Baltis Vallis, is ~6,800 km long. The large width of most channels implies that they are not simply collapsed lava tubes. The channels have no tributaries and mostly occur in association either with the lower subunit of regional plains or with lobate plains. The length-frequency distribution of the channels displays two modes corresponding to shorter (<~150 km) and longer (>~150 km) features. The longer channels usually occur within the areas of the lower subunit of regional plains, whereas the shorter channels are associated predominantly with fields of lobate plains.

The great length of the channels requires extremely fluid materials for their formation. Potential channel-forming fluids with low viscosity (about 0.1 Pa s) include ultramafic komatiites, high-Ti lunar-type basalts, carbonatites, and sulfur flows. Carbonatites and sulfur flows are attractive alternatives, especially due to the abundance of sources for carbon and sulfur on Venus. One of the likely mechanisms of channel formation is related to the dense atmosphere of Venus, which rapidly cools lavas to form discontinuous, plate-like crust over fluid filling the channels. Such a roof thermally insulates the fluid and enables it to remain molten over great distances.

Volcano-Tectonic Features

This category includes coronae, arachnoids, novae, and calderas. These features consist of a volcanic component in form of lava flows and a tectonic component that defines the shape of the features.

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Figure 9. Examples of volcano-tectonic features on Venus. (a) Coronae. Branches of groove belts (gb) form the circular rims of the coronae; center of the image is at 20.8°S, 220.2°E. (b) Arachnoids. These circular and elliptical features (arrows) are surrounded by topographic rims that, at some arachnoids, are cut by swarms of concentric fractures; center of the image is at 40.0°N, 19.2°E. (c) A nova. The star-like pattern of grooves defines the nova, which is completely inside the rim of Pavlova Corona (radar-bright circular feature). Some of the narrow graben radiating from the nova cut both the rim and the embaying plains (white arrows). To the east of the center of the nova, flows of lobate plains appear to emanate from the graben of the nova (black arrow); center of the image is at 14.5°N, 39.0°E. (d) Calderas (white arrows). These are circular topographic depressions outlined by swarms of concentric fractures. Some of the calderas appear as the sources of lobate plains (black arrow); center of the image is at 30.5°S, 224.5°E.

Coronae (Figure 9a) are circular structures consisting of an annulus (rim) of fractures (rarely, ridges) that surrounds interiors with numerous volcanic and tectonic landforms. There are about 200 coronae on Venus. Structures of the rim represent the tectonic component of coronae. Lava flows radiating away from the rim and sometimes partly filling the corona interiors represent the volcanic component. The total diameter range of the corona rims is from ~70 up to ~1,000 km, with a peak at ~200–250 km.

Because of their circular shape and large diameters, coronae are interpreted as the surface manifestations of mantle diapirs. The predominance of fractures and graben in corona rims suggests that coronae formed in extensional environments, which may be related to both the topographic evolution of the surface above the diapir and the intrusion of circular dikes.

The volcanic components of coronae postdate emplacement of the structures that form corona rims. The most common volcanic features that are associated with coronae are shield plains (~36% of coronae) and lobate/smooth plains (~35%). For ~22% of coronae, the latest volcanic activity is in the form of flows of the upper subunit of regional plains. Thirty-seven coronae (~17%) are flooded by the lower subunit of regional plains and do not show evidence of volcanism.

Arachnoids (Figure 9b) constitute another type of volcano-tectonic feature. These features usually have a broad and apparently low topographic rim, with or without fractures, and are surrounded by radial wrinkle ridges. About 250 arachnoids have been detected on Venus; they are predominantly associated with the upper subunit of regional plains and control the distribution of wrinkle ridges in their close vicinity (Figure 9b).

Novae (Figure 9c): Abundant tectonic structures characterize novae, and the volcanic components of these volcano-tectonic features are either subdued or absent. When volcanic features are associated with novae, they represent flows of lobate plains that emanate from the graben of novae. This suggests that the graben are the surface manifestations of dikes, and novae may represent a specific class of radiating dike swarms.

Calderas (Figure 9d): About 100 calderas are listed in the catalogue of volcanic features of Venus (Crumpler & Aubele, 2000). Calderas represent broad (60–80 km) and shallow (hundreds of meters) topographic depressions surrounded by swarms of concentric fractures. Calderas are associated with shield plains (~31% of the features), the upper subunit of regional plains (~41%), and lobate plains (~21%). About 7% of calderas occur within fields of the lower subunit of regional plains.

Tectonic Landforms

Tectonic landforms on Venus are represented by terrains in which tectonic structures play a dominant role and define the final morphology of the terrain. Tectonic structures are always secondary features relative to material(s) that they cut. In some cases, however, the structures are strongly coupled with the material component of the terrain and define material-structural units. In the other cases, however, tectonic structures run through a suite of material complexes and form purely structural units.

The Surface of VenusClick to view larger

Figure 10. Examples of tectonic landforms on Venus. (a) Tessera terrain is defined by numerous chaotically oriented structures of contractional and extensional origin. Tessera represents one of the most heavily tectonized units on Venus; center of the image is at 69.9°N, 18.7°E. (b) Fine-scale lineaments dissect the surface of densely lineated plains (pdl). Materials of the plains embay a piece of tessera (lower left, t); lineaments of pdl penetrate into tessera (arrow). These relationships indicate that both emplacement and deformation of pdl postdate tessera; center of the image is at 31.9°N, 128.0°E. (c) Sub-parallel, curvilinear ridges deform the surface of ridged plains (pr/RB). In places, the ridges are collected in prominent ridge belts (lower right). Near the contacts with tessera (t), materials of the plains embay pieces of tessera and structures of the plains/belts cut tessera; center of the image is at 1.9°N, 111.2°E.

Tessera terrain (t, Fortuna Formation, Figure 10a): Tessera is a material-structural unit in which both the material and structural components occur together within large massifs. The exposed surface of tessera, ~35.7 × 106 km2, composes about 7% of the surface of the planet (Table 4). Tessera massifs vary in sizes from a few tens up to a few thousands of kilometers and tend to cluster within several large regions (Figure 6a). The largest such region in the equatorial zone of Venus is about 5,000 by 14,000 km. The largest tesserae in this region (Ovda and Thetis Tesserae) correspond to a plateau-like class of regional highlands on Venus.

At least two sets of intersecting structures of contractional (ridges) and extensional (graben, fractures, scarps) origin characterize the surface of tessera. These structures vary in size from several kilometers up to a few hundred kilometers. Broad (tens of kilometers wide) and very long (many hundreds of kilometers) troughs divide the largest tessera regions into individual segments. Intratessera lava plains often partly cover the floor of these troughs.

The age relationships between the contractional and extensional structures play a key role in the assessment of the models of tessera formation. These relationships are mostly unclear and cannot be established unambiguously because of the lack of stratigraphic markers. In some places, however, the intratessera plains provide such markers that intervene between the ridges and graben and represent a robust indicator of the stratigraphic relationships of the tessera structures. In all tessera regions where the stratigraphic relationships between the tessera ridges, graben, and intratessera plains are clear, the plains indicate the older age of the ridges.

Densely lineated plains (pdl, Atropos Formation, Figure 10b): Numerous densely packed lineaments, narrow, short, and parallel/subparallel to each other, dissect the surface of this material-structural unit. If the lineaments are wide enough, they usually appear as fractures. Commonly, the densely packed lineaments completely erase the morphology of the precursor materials. In some occurrences of the unit, however, remnants of preexisting lava plains are visible between the lineaments. Densely lineated plains occupy a small area of Venus, about 7.2 × 106 km2 (1.6% of the planet surface, Table 4), and are observed as slightly elevated and usually small (tens of kilometers across) patches.

Ridged plains (pr, Lavinia Formation, Figure 10c): Ridged plains represent another material-structural unit, in which tectonic structures, broad (5–10 km wide) and long (several tens of kilometers) ridges, deform the material component presented by volcanic plains. Very often, the ridges are collected into prominent and elevated belts (ridge belts) that are many hundreds of kilometers long and many tens of kilometers wide. Ridged plains and ridge belts occupy about 9.6 × 106 km2 (2.1% of the surface of Venus, Table 4) and occur in many regions on Venus (Figure 6a).

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Figure 11. Examples of groove belts and rift zones. (a) Swarms of narrow curvilinear graben form groove belts (gb). The graben cut underlying materials and eliminate most of their morphologic characteristics; center of the image is at 40.8°S, 341.7°E. (b) Rift zones (rz) also consist of multiple extensional structures that are wider and longer than the grooves of groove belts; center of the image is at 9.4°N, 204.6°E.

Groove belts (gb, Agrona Formation, Figure 11a): Groove belts are defined as zones of densely packed extensional structures, fractures and graben. The typical widths of these structures are from several hundred meters up to 1–2 km, and individual fractures can reach several tens of kilometers in length. The density of structures in groove belts is usually so high that they almost completely obscure the morphology of underlying materials. Groove belts represent a purely tectonic unit during the formation of which various materials have been cut by extensional structures; they occupy ~8.7% of the surface of Venus (Table 4).

Groove belts occur as zones of many hundreds (up to a few thousands) of kilometers long and tens to a few hundred kilometers wide that are distributed more broadly than ridge belts (Figure 6a) and form an anastomosing pattern of branches. Fractures and graben of the branches often form the rims of coronae. Branches of groove belts always represent local highs, the surface of which stands several hundred meters above the surrounding plains.

Rift zones (rz, Devana Formation, Figure 11b): Rift zones, as well as groove belts, represent a purely structural type of terrain that consists of numerous and densely packed extensional structures. In many rift zones, there are deep (several kilometers) and steep-sided canyons that can be several tens of kilometers wide. In contrast to groove belts, structures of rift zones on average are broader, longer, and somewhat less densely packed. Fractures and graben of rift zones are either parallel to each other or have a zigzag-like planform, which is not typical of structures of groove belts. Rifts appear as broad (hundreds of kilometers wide) and very prominent zones that extend for thousands of kilometers. They occupy ~5.0% of the surface of Venus (Table 4), preferentially occur within the equatorial region of Venus, and outline the BAT (Beta-Atla-Themis) region (Figure 6a).

Age Relationships of the Main Volcanic and Tectonic Landforms

The volcanic and tectonic landforms of Venus make up about 99.5% of its surface (impact craters occupy the rest, ~0.5%) and, in fact, represent pieces of the geological record of the planet. Thus, the understanding of the relative age relationships among the volcanic and tectonic terrains is crucially important to unraveling the geological history of Venus. The analysis of the embayment and crosscutting relationships is used to establish the relative ages among the terrains at the local scale (Basilevsky & Head, 1995). At the next stage, continuous geological mapping is required to extend the local stratigraphy first to the regional and then to the global scales. Geological mapping of the majority (~97%) of the surface of Venus (Ivanov & Head, 2011) allowed detail documentation and analysis of the boundaries and the embayment/crosscutting relationships between the terrain virtually in all regions of Venus. As a result, the global-scale stratigraphic scheme of Venus was constructed (Figure 6b).

Practically everywhere on Venus, contacts of tessera with surrounding plains are very sinuous due to penetration of materials of plains into tessera massifs (Figure 12a). All tectonic structures of tessera are abruptly truncated by the contact with the plains. These relationships provide compelling evidence that both the emplacement of tessera materials and their tectonic modification were completed before formation of the younger vast volcanic plains. The same relationships with the surrounding plains characterize occurrences of both densely lineated plains and ridged plains/ridge belts (Figure 12b, c). Occurrences of these tectonized terrains are embayed by material of the adjacent plains, which thus are younger. Structures of groove belts usually cut tessera, densely lineated plains and ridged plains/ridge belts, which indicate that groove belts formed later than these terrains. The abundance of groove belts on Venus (Figure 6a) allows tracing these relationships in many areas, and everywhere the belts appear to be younger. Vast plains units such as shield plains and both subunits of regional plains embay groove belts (Figure 12d).

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Figure 12. Relative age relationships between the older tectonized units and the vast volcanic plains (psh and rp1). The surface of the vast volcanic plains is mildly deformed, and materials of the plains completely superpose all structures of the tectonized units. These relationships indicate the younger age of the plains. (a) Center of the image is at 47.3°N, 127.6°E. (b) Center of the image is at 40.9°N, 34.7°E. (c) Center of the image is at 37.5°N, 156.4°E. (d) Center of the image is at 3.0°N, 146.7°E.

Thus, the absolute majority of the tectonically deformed terrains on Venus (Figure 6b) predate emplacement of both shield and regional plains. These older tectonized terrains make up about 20% of the surface of Venus, and this is the minimum estimate of their abundance, because the younger volcanic plains certainly have buried some portion of them.

Fractures and graben of rift zones cut all tectonic units, shield plains, and regional plains (Figure 13a) and thus are younger (Figure 6b). Rift zones occur in close spatial association with lobate plains (Figure 6a). Structures of rifts both cut the plains and are embayed by their material (Figure 13b), which indicates contemporaneous formation of the rifts and lobate plains.

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Figure 13. Relative age relationships of rift zones with the adjacent units. (a) When rift zones (rz) cross regional plains (rp1), structures of rifts cut the surface of the plains and thus postdate this unit and all older units; center of the image is at 12.2°N, 199.3°E. (b) Rift zones and lobate plains (pl) show relationships that suggest their broadly contemporaneous formation: in some places structures of rifts cut the plains (black arrows), and in a neighboring region lobate plains embay structures of rift zones (white arrows); center of the image is at 6.2°N, 199.8°E. (From Ivanov & Head, 2013, with changes.)

The main volcanic units on Venus also demonstrate consistent relationships of relative ages among each other at the global scale. The most abundant volcanic units are shield plains and the lower subunit of regional plains (Table 4). Owing to their abundance, contacts between these units are very frequent and can be studied in detail in many regions of Venus. Three hypothetical types of the age relationships between these units can be envisioned: (1) shield plains are older (Figure 14a), shield clusters and regional plains are quasi-synchronous (Figure 14b), and shield clusters are younger (Figure 14c) than the adjacent regional plains. Each type of the relative age relationships defines a set of criteria that can be applied to a real situation on the surface of Venus in order to establish the relative aged between the shield and regional plains.

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Figure 14. A sketch illustrating three possible age relationships of small shields and regional plains: (a) shields are older, (b) shields are contemporaneous, and (c) shields are younger than regional plains. The shields and intershield plains from the first example were mapped as shield plains, and the other shields were lumped into a unit of shield clusters in the global geological map. (From Ivanov & Head, 2013, with changes.)

Application of these criteria to randomly selected occurrences of shield plains has shown that in ~70% of analyzed cases regional plains are younger. About 10% of analyzed shield fields appear to be synchronous with regional plains, and ~10% of shield clusters postdate adjacent regional plains. The rest of the analyzed population shows either ambiguous or obscured age relationships with regional plains.

In regions where both subunits of regional plains show clean and sharp contact, it is seen that material of the upper subunit either penetrates into local lows of the lower subunit or embay local highs of the lower subunit (Figure 7b). These relationships are observed globally and indicate the relatively young age of unit rp2.

In all places where lobate plains occur in contact with either shield plains or any of the subunits of regional plains, flows of lobate plains are younger.

Geologic History of Venus

The globally observed stratigraphic relationships among the tectonic and volcanic units on Venus divide the observable portion of its geologic history into three different episodes, each with a specific style of resurfacing. These are as follows (Figure 15): (1) Global tectonic regime, when tectonic resurfacing dominated. Exposed occurrences of these units comprise about 20% of the surface of Venus (Table 4). (2) Global volcanic regime, when volcanism was the most important process of resurfacing and resurfaced about 60% of Venus (Table 4). (3) Network rifting-volcanism regime, when both tectonic and volcanic activity were about equally important. During this regime, about 16% of the surface of Venus was modified (Table 4).

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Figure 15. A global correlation chart that shows the three major regimes of resurfacing on Venus. (From Ivanov & Head, 2015, with changes.)

Global Tectonic Regime

Relationships of relative ages between the ridges and graben in tessera are crucially important for constraining the possible models of tessera formation and thus the global tectonic style near the beginning of the recognizable history of Venus. In the proposed upwelling or the lava pond models of tessera formation, the extensional structures (graben) are considered as the oldest structures formed by stretching and cracking of either the roof of the rising molten diapir or a large lava pond. In contrast, in the downwelling models thickening of the crust over the sites of mantle down flow due to under- and overthrusting, bending, and buckling of crustal slabs has resulted in the initial formation of broad contractional structures (ridges).

The robust evidence for the older age of the tessera ridges (Figure 16) suggests that individual tesserae more likely represent large and approximately equidimensional sites of primary contractional structures. Regions in which tessera massifs cluster (Figure 6a) may represent the loci of downwelling. The dimensions of these regions (several thousands of kilometers) may correspond to the upper limit of the sizes of the downwelling cells. The major structural seams (e.g., long narrow troughs that divide tesserae into series of blocks) characterize individual tessera massifs but do not extend from one large tessera into the other. This suggests that individual large tesserae evolved independently, and their typical dimensions (many hundreds to a few thousands of kilometers) may characterize the lower limit of the sizes of the downwelling cells.

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Figure 16. An example of stratigraphic relationships of contractional and extensional structures in tessera where lava plains serve as a local stratigraphic marker that indicates the relative age of ridges and graben. The figure shows a piece of the surface in the northern-central portion of Ovda tessera. The older intratessera plains (plit1) embay the tessera ridges and are cut by sets of narrow graben; center of the image is at 2.3°S, 90.0°E. (From Ivanov & Head, 2015, with changes.)

Ridge belts extend for many hundreds of kilometers and consist of contractional structures. This suggests that the belts formed under compressional stresses applied within relatively narrow but very long zones. These characteristics resemble those of terrestrial thrust-and-fold belts, the structures of which are interpreted to form over large thrust faults and indicate crustal shortening due to lateral movements of crust/lithosphere. If similar processes participated in formation of ridge belts on Venus, then the relatively low relief of the belts (a few hundred meters) suggests that the lateral movements and related contraction were rather small.

The distinctive exceptions to this are the mountain belts around Lakshmi Planum. The belts represent the highest mountain ranges on Venus, and their relationships with the surrounding terrains provide evidence for large-scale shortening, collision, underthrusting, and epeirogenic uplift. The mountain belts, however, exist only in one region, which suggests that even if the belts are related to processes akin to subduction, they were fairly restricted on Venus.

Thus, the earlier phases of the global tectonic regime (Figure 15) were characterized by the predominance of contractional structures (tessera ridges, ridge belts, mountain belts) that were likely to have been related to lateral, but limited, movements of the lithosphere.

In comparison with the other terrains of the global tectonic regime, groove belts are more broadly distributed over the surface and do not show evidence of association with the tessera clusters (Figure 6a). The consistent relative age relationships of groove belts imply that the belts formed during the later phases of the global tectonic regime (Figure 15). The fractures and graben that define the belts imply that these phases were predominantly related to broadly distributed extension of the crust/lithosphere. An important characteristic of groove belts is that their branches often represent the tectonic components of coronae (Figure 9a) that are thought to be the surface manifestations of mantle diapirs. Thus, the close spatial association of the belts and coronae suggest that these features formed mutually due to multiple and relatively small-scale mantle upwellings/diapirs. The tight spatial association of groove belts with coronae may indicate that the dominant style of resurfacing at the later phases of the global tectonic regime was plume tectonics that caused the predominantly vertical displacement of lithosphere and its deformation.

Global Volcanic Regime

The earlier tectonic regime was followed by emplacement of the vast volcanic plains (shield and regional plains). Emplacement of the plains defines the second, globally volcanic regime, when tectonic deformation related to mantle convection waned and volcanic resurfacing dominated (Figure 15). The main units of the global volcanic regime have an obviously different morphology, and this is an indication of different volcanic styles during their formation.

The most obvious features of shield plains are small and very abundant volcanic constructs (Figure 7a). The great abundance of the constructs implies that their sources were fairly pervasive and nearly globally distributed, while the small sizes of the shields suggest that the supply of magma in their sources was limited. Another important feature of shield plains is that the steep-sided domes are spatially and stratigraphically associated with them. These associations favor the higher silica content in the parent magma of the domes, for example, because of partial remelting of basaltic crust. The small size of the constructs of shield plains and their association with the steep-sided domes are most consistent with shallow crustal melting and differentiation of magma in reservoirs and/or partial melting of the crustal materials.

The lower subunit of regional plains postdates emplacement of shield plains and forms very broad and morphologically homogenous surfaces that are distinctly different morphologically from the preceding shield plains (Figure 7b). Two important features characterize this unit and provide the keys for the understanding of its mode of formation. (1) The sources of lavas are not visible at the resolution of Magellan images. (2) The lower subunit of regional plains are very abundant and ubiquitous: the exposed surface of this unit comprises about one-third of the surface of Venus (Table 4), and the unit occurs almost everywhere on the planet (Figure 6a). These characteristics of the plains strongly suggest that they formed by voluminous volcanic eruptions from broadly distributed sources, during which individual lava flows coalesced into an essentially single volcanic unit and buried the source regions. The absence of noticeable volcanic constructs suggests that eruptions of materials of the plains were extremely voluminous and relatively short-lived.

Large and radar-bright flows characterize the upper subunit of regional plains (rp2). Usually, the flows of this unit clearly indicate the source areas of the plains and are represented by individual features such as large and intermediate volcanoes and some coronae. The distinct volcanic flows of the upper subunit of regional plains and their clear association with individual sources mark the other change of volcanic style on Venus, from the massive and broadly distributed eruptions to more localized volcanic activity at fewer specific centers.

The density of impact craters on terrains of the global tectonic regime (t, pdl, pr/RB, and gb) and global volcanic regime (psh, rp1, rp2) are practically indistinguishable (Figure 17). The stratigraphic relationships between these terrains, however, are always clear and unambiguously indicate the younger age of the plains of the volcanic regime. This suggests a rapid change from the earlier preferentially tectonic to the following essentially volcanic resurfacing.

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Figure 17. Mean crater densities on units that formed during different regimes of resurfacing on Venus. The densities for the global tectonic and the global volcanic regimes are statistically indistinguishable, but the crater density on rift zones and lobate plains (the network rifting-volcanic regime) is much smaller. Note that error bars of the density estimates are four sigma. (From Ivanov & Head, 2015, with changes.)

Network Rifting-Volcanism Regime

The network rifting-volcanism regime characterized the later episodes of the geologic history of Venus. Rift zones represent a tectonic component of this regime, and lobate plains represent its volcanic component (Figure 15).

Rift zones are very prominent morphologically, but their total area is about four times smaller than the exposed (i.e., minimal) area of tectonic terrains of the global tectonic regime (Table 4). In contrast to these earlier tectonic terrains, which are broadly distributed over the surface of Venus, rift zones form a few very long and pronounced zones that are prominently concentrated in the equatorial and the BAT regions of Venus (Figure 6a). These characteristics of rift zones suggest that (1) the effect of tectonic resurfacing diminished with time throughout the visible geologic history and (2) the style of tectonic resurfacing evolved from the earlier broadly distributed deformation to the highly concentrated deformation during the later network rifting-volcanism regime. These changes in both the intensity and the lateral extent of tectonic resurfacing are consistent with the transition from an earlier mobile-lid regime of mantle circulation with thinner/weaker lithosphere to a later stagnant-lid regime with thicker/stronger lithosphere (e.g., Parmentier & Hess, 1992; Head et al., 1994; Solomatov & Moresi, 1996; Grimm & Hess, 1997).

Another possible line of evidence for the thicker thermal lithosphere during the network rifting-volcanism regime is that rift zones are poorly correlated with coronae but are clearly associated with the large dome-shaped rises of the BAT region. This correlation suggests that during this time the large-scale mantle upwellings/plumes (likely manifested by the rises) controlled the spatial distribution of the principal zones of extension and rupture of the crust. The smaller-scale mantle diapirs (likely manifested by coronae) appear to have played a subordinate role in resurfacing during the network rifting-volcanism regime. This may be because the thicker/stronger lithosphere at this time acted as a rheological barrier that filtered out smaller mantle diapirs but was more “transparent” to the larger mantle upwellings/plumes.

The mean density of craters superposed on the surface of both rift zones and lobate plains is significantly lower than the crater densities on units of the global tectonic and volcanic regimes (Figure 17). This suggests that the transition from the global volcanic regime to the fully developed network rifting-volcanism regime was longer than the change from the global tectonic to global volcanic regimes. The apparently abrupt change from tectonically to volcanically dominated styles of resurfacing may correspond to a rapid change in mantle convection, whereas the extended transition from the global volcanic to the network rifting-volcanism regimes may reflect a gradual increase of the lithosphere thickness.

In conclusion, Earth-based and orbital observations of the surface of Venus have significantly advanced our understanding of the geological evolution of this planet. The major diversity between Earth and Venus, their global tectonic styles, appears to emerge. Plate tectonics that has continuously operated on Earth for billions of years is in sharp contrast with the apparently time-dependent regimes of resurfacing on Venus. Thus, both planets likely follow their own, dissimilar, ways of evolution. What are the principal factors of these differences? Studies of the surface morphology and relationships of landforms on the surface are not enough to answer this question, and progress in understanding of the geochemical history of the planet is needed as well.

There are still very limited data on Venus geochemistry, which prevents the assessment of its evolution from different points of view. The necessary geochemical data can be obtained only by the direct measurement of the chemical composition of soils on the surface. Thus, new lander-oriented missions to Venus are required to close a part of the large gap in our geochemical knowledge of Venus. What is the composition of each of the major varieties of volcanic plains? How do they compare with volcanic products from different geodynamic environments on Earth? Chemical in situ measurements will provide important constraints on these problems.

One of the major pieces of the Venus geological and geochemical puzzle is the composition of tessera material. Tessera is one of the oldest terrains on Venus and likely represents a window into its geological past. Tessera material may bear geochemical markers of the earlier epochs of the evolution of the planet and is therefore a target of the highest scientific interest. Another problem of the geochemistry of Venus is related to the presence or absence of non-basaltic material on this planet. The steep-sided domes and festoons currently provide the only morphologic evidence of such material. Both the abundance and dimensions of these volcanic constructs suggest that a non-basaltic component may compose a noticeable fraction of the Venusian crust. Again, new data on geochemistry are needed to assess the possible compositional diversity of the crust on Venus.

Despite significant progress in our understanding of the surface geology of Venus, several important issues remain.

As a large terrestrial planet, Venus may still be active volcanically and tectonically. The recent ESA Venus Express mission provided some circumstantial evidence for possible current volcanic activity on the planet. In order to test these findings, it is necessary to conduct a long-term monitoring of the surface from orbit with the help of high-resolution radar systems. In addition, orbital high-resolution observations of the surface will help to address the other major problem of Venus geology, which is related to the possible mode of the emplacement of the lower subunit of regional plains. These plains are the most abundant on Venus and thus play an important role in the understanding of the thermal budget of the planet. In contrast to other plains units, the lower subunit of regional plains does not display its source areas at the Magellan resolution. Does this unit represent the result of fissure eruptions, or was it emplaced from distributed small edifices? Is it a true single volcanic unit, as is observed at the available resolution, or does it consist of a number of coalesced individual volcanic fields? The answers to these questions are very important for the formulation of adequate models of the geologic evolution of Venus.

The new missions to Venus that will include both landers and long-term orbiters for surface observations are needed to solve the vital outstanding problems of Venus geology.

Further Reading

Basilevsky, A. T., & Head, J. W. (1998). The geologic history of Venus: A stratigraphic view. Journal of Geophysical Research, 103, 8531–8544.Find this resource:

Ford, P. G., & Pettengill, G. H. (1992). Venus topography at kilometer-scale slopes. Journal of Geophysical Research, 97, 13103–13114.Find this resource:

Grimm, R. E. (1994). The deep structure of Venusian plateau highlands. Icarus, 112, 89–103.Find this resource:

Hauck, S. A., Phillips, R. J., & Price, M. H. (1998). Venus: Crater distribution and plains resurfacing models. Journal of Geophysical Research, 103, 13635–13642.Find this resource:

Head, J. W., Crumpler, L. S., Aubele, J. C., Guest, J. E., & Saunders, R. S. (1992). Venus volcanism: Classification of volcanic features and structures, associations, and global distribution from Magellan data. Journal of Geophysical Research, 97, 13153–13197.Find this resource:

Ivanov, M. A., & Head, J. W. (2013). The history of volcanism on Venus. Planetary and Space Science, 84, 66–92.Find this resource:

Ivanov, M. A., & Head, J. W. (2015). The history of tectonism on Venus: A stratigraphic analysis. Planetary and Space Science, 113–114, 10–32.Find this resource:

Ivanov, M. A., & Head, J. W. (2015). Volcanically embayed craters on Venus: Testing the catastrophic and equilibrium resurfacing models. Planetary and Space Science, 106, 116–121.Find this resource:

Phillips, R. J., Raubertas, R. F., Arvidson, R. E., Sarkar, I. C., Herrick, R. R., Izenberg, N., & Grimm, R. E. (1992). Impact craters and Venus resurfacing history. Journal of Geophysical Research, 97, 15923–15948.Find this resource:

Schaber, G. G., Strom, R. G., Moore, H. J., Soderblom, L. A., Kirk, R. L., Chadwick, D. J., . . . Russel, J. (1992). Geology and distribution of impact craters on Venus: What are they telling us? Journal of Geophysical Research, 97, 13257–13301.Find this resource:

Schubert, G., Solomatov, V. S., Tackley, P. J., & Turcotte, D. L. (1997). Mantle convection and the thermal evolution of Venus. In S. W. Bougher, D. M. Hunten, & R. J. Phillips (Eds.), Venus II: Geology, geophysics, atmosphere, and solar wind environment (pp. 1245–1287). Tucson: University of Arizona Press.Find this resource:

Solomatov, V. S. (2004). Initiation of subduction by small-scale convection. Journal of Geophysical Research, 109, B01412.Find this resource:

Solomon, S. C., Smrekar, S. E., Bindschadler, D. L., Grimm, R. E., Kaula, W. M., McGill, G. E., . . . Stofan, E. R. (1992). Venus tectonics: An overview of Magellan observations. Journal of Geophysical Research, 97, 13199–13255.Find this resource:

References

Abdrakhimov, A. M. (2005). Geochemical comparison of volcanic rocks from terrestrial intraplate oceanic hot spots with Venusian surface material. Geochemistry International, 43, 732–747.Find this resource:

Barsukov, V. L., Basilevsky, A. T., Burba, G. A., Bobinna, N. N., Kryuchkov, V. P., Kuzmin, R. O., . . . Akim, E. L. (1986). The geology and geomorphology of the Venus surface as revealed by the radar images obtained by Venera 15 and 16. Journal of Geophysical Research, 91, D399–D411.Find this resource:

Basilevsky, A. T., & Head, J. W. (1995). Regional and global stratigraphy of Venus: A preliminary assessment and implications for the geological history of Venus. Planetary and Space Science, 43, 1523–1553.Find this resource:

Campbell, D. B., & Burns, B. A. (1980). Earth-based radar imagery of Venus. Journal of Geophysical Research, 85, 8271–8281.Find this resource:

Crumpler, L. S., & Aubele, J. (2000). Volcanism on Venus. In B. Houghton, H. Rymer, J. Stix, S. McNutt, & H. Sigurdson (Eds.), Encyclopedia of volcanoes (pp. 727–770). San Diego: Academic Press.Find this resource:

Grimm, R. E., & Hess, P. C. (1997). The crust of Venus. In S. W. Bougher, D. M. Hunten, & R. J. Phillips (Eds.), Venus II: Geology, geophysics, atmosphere, and solar wind environment (pp. 1163–1204). Tucson: University of Arizona Press.Find this resource:

Head, J. W., Crumpler, L. S., Aubele, J. C., Guest, J. E., & Saunders, R. S. (1992). Venus volcanism: Classification of volcanic features and structures, associations, and global distribution from Magellan data. Journal of Geophysical Research, 97, 13153–13197.Find this resource:

Head, J. W., Parmentier, E. M., & Hess, P. C. (1994). Venus: Vertical accretion of crust and depleted mantle and implications for geological history and processes. Planetary and Space Science, 42, 803–811.Find this resource:

Ivanov, M. A., & Head, J. W. (2011). Global geological map of Venus. Planetary and Space Science, 59, 1559–1600.Find this resource:

Parmentier, E. M., & Hess, P. C. (1992). Chemical differentiation of a convecting planetary interior: Consequences for a one plate planet such as Venus. Geophysical Research Letters, 19, 2015–2018.Find this resource:

Saunders, R. S., Spear, A. J., Allin, R. S., Austin, R. S., Berman, A. L., Chandlee, R. C., . . . Wall, S. D. (1992). Magellan mission summary. Journal of Geophysical Research, 97, 13067–13090.Find this resource:

Solomatov, S. V., & Moresi, L.-N. (1996). Stagnant lid convection on Venus. Journal of Geophysical Research, 101, 4737–4753.Find this resource:

Solomon, S. C., Smrekar, S. E., Bindschadler, D. L., Grimm, R. E., Kaula, W. M., McGill, G. E., . . . Stofan, E. R. (1992). Venus tectonics: An overview of Magellan observations. Journal of Geophysical Research, 97, 13199–13255.Find this resource:

Surkov, Y. A. (1997). Exploration of terrestrial planets from spacecraft: Instrumentation, investigation, interpretation (2nd ed.). Chichester: Praxis.Find this resource: