identification of mercurian volcanism: resolution effects and ...to be composed of ejecta, rather...

Meteoritics & Planetary Science 37, 1209–1222 (2002)Available online at http://www.uark.edu/meteor

© Meteoritical Society, 2002. Printed in USA.1209

Prelude preprint MS#4734

Identification of mercurian volcanism: Resolution effects and implicationsfor MESSENGER

S. M. MILKOVICH1*, J. W. HEAD1 AND L. WILSON2

1Department of Geological Sciences, Brown University, Providence, Rhode Island 02912, USA2Department of Environmental Science, Lancaster University, Lancaster LA1 4YQ, U.K.

*Correspondence author's e-mail address: [email protected]

(Received 2002 January 15; accepted in revised form 2002 May 17)(Presented at the Workshop on Mercury, The Field Museum, Chicago, Illinois, 2001 October 4–5)

Abstract–The possibility of volcanism on Mercury has been a topic of discussion since Mariner 10returned images of half the planet's surface showing widespread plains material. These plains couldbe volcanic or lobate crater ejecta. An assessment of the mechanics of the ascent and eruption ofmagma shows that it is possible to have widespread volcanism, no volcanism on the surface whatsoever,or some range in between. It is difficult to distinguish between a lava flow and lobate crater ejectabased on morphology and morphometry. No definite volcanic features have been identified on Mercury.However, known lunar volcanic features cannot be identified in images with similar resolutions andviewing geometries as the Mariner 10 dataset. Examination of high-resolution, low Sun angle Mariner10 images reveals several features which are interpreted to be flow fronts; it is unclear if these arevolcanic flows or ejecta flows. This analysis implies that a clear assessment of volcanism on Mercurymust wait for better data. MESSENGER (MErcury: Surface, Space ENvironment, GEochemistry,Ranging) will take images with viewing geometries and resolutions appropriate for the identificationof such features.

INTRODUCTION

Ever since the arrival of Mariner 10 at Mercury in 1974–1975there has been debate over the existence of volcanic featureson the surface of the planet. Mercury has two types of plains,smooth and intercrater; they are the most extensive terrain onthe planet (Strom et al., 1975). The smooth plains arecharacterized by flat or gently rolling topography overlain bywrinkle ridges. The intercrater plains are level to gently rollingas well, and are found between and around the cratered terrainthat covers much of the mercurian surface (Spudis and Guest,1988). Through comparisons with lunar mare, manyresearchers (e.g., Murray et al., 1975; Strom et al., 1975)concluded that these plains were formed by volcanic processes;the smooth plains look very similar to the lunar mare and arequite widespread. Several plains fill up basins and craters whileothers sit in low-lying areas between craters. The widespreaddistribution of the intercrater plains indicates extensiveresurfacing early in Mercury's history and many researchersconclude that the morphology of the plains is most consistentwith volcanism as a resurfacing agent (Murray et al., 1975;Strom et al., 1975; Dzurisin, 1978; Kieffer and Murray, 1987).Others consider that they are more similar to some of the lunarlight plains formed by crater ejecta (Wilhelms, 1976). Wilhelms

pointed out that the pro-volcanism arguments are similar toarguments used to demonstrate that certain light plains (theCayley Formation) on the Moon were volcanic; Apollo 16samples proved these plains to be made up of impact breccia.The plains may have been formed by fluidized ejecta, impactmelt, overlapping ejecta deposits from multiple craters, orsecondary crater ejecta. This debate can only be resolved iflandforms typical of extrusive volcanism (e.g., domes, shields,sinuous rilles, calderas, flow fronts) are clearly identified orclearly found not to exist on the surface of Mercury.

Schultz (1977) pointed out that lunar volcanic features aresmall, localized, and nearly undetectable at resolutionscomparable to those of Mariner 10. Work by Malin (1978)indicated that the viewing geometry of Mariner 10 may havebeen insufficient for the identification of such volcanic features.He addressed this issue by analyzing lunar images underMariner 10 resolutions (∼ 1–2 km) and various lightingconditions (0–30°) for volcanic features. Since the mercurianimages are at about the same resolution as pre-1964 Earth-based telescopic images of the Moon, Malin used such imagesfrom the Consolidated Lunar Atlas (Kuiper et al., 1967) withapproximately the same viewing conditions as the Mariner 10data. Malin was only able to identify lunar volcanic domesand flow fronts under these conditions, and found that volcanic

1210 Milkovich et al.

features cannot be identified in images with Sun elevationangles above 25°. Only 15% of the surface was imaged withSun elevation angles <20° (Malin, 1978). Smooth plains wereidentified in 40% of the images, or 18% of the surface (Spudisand Guest, 1988). Due to the limited coverage of mercurianplains at low Sun angles, Malin was not surprised that volcaniclandforms have not been identified on Mercury. He concludedthat the only way to identify volcanic landforms in the Mariner 10data is to look near the terminator. In such images he identifiedtwo possible volcanic domes on Mercury; these are positivetopographic features which are reminiscent of lunar domes butare near scarps and other structural features and are notunequivocally volcanic. One of these possible domes is locatednear Discovery Scarp, and can be seen in Fig. 1. Recentrecalibration of Mariner 10 color data has allowed identificationof several features spectrally similar to volcanic flows andpyroclastic fissure eruptions (Robinson and Lucey, 1997). Therecalibrated data (Robinson and Taylor, 2001) and microwaveand mid-infrared observations (Jeanloz et al., 1995) indicate alow abundance of FeO in the crust of Mercury; this may implythat basalt is absent on the mercurian surface (Sprague et al.,1994; Jeanloz et al., 1995) or that the mantle has a low FeOcontent (Robinson and Taylor, 2001).

The question of the mode of formation of the mercurianplains remains unanswered. This study assesses the bestway of identifying volcanic landforms on Mercury byaddressing several key issues. These issues are whatdistinguishes an ejecta flow from a volcanic plain, whatwould mercurian volcanism look like, how would we findevidence for mercurian volcanism in images, and what dowe actually see in the Mariner 10 images. Finally, we discussthe upcoming MESSENGER (MErcury, Surface, SpaceENvironment, GEochemistry, Ranging) spacecraft missionto Mercury and examine its prospects for identifying possiblevolcanic landforms.

LOBATE EJECTA FLOWS

Outline of Process

Early analyses of lunar ejecta deposits relied heavily onterrestrial analogs, particularly small impact craters such asMeteor Crater, Arizona, and craters resulting from high-explosive and nuclear events. The small sizes of these meantthat the ejecta velocities were such that most near-crater ejectawas emplaced as an "overturned flap" and at relatively lowvelocities. Consideration of the full temporal and spatial rangeof planetary craters, however, soon showed that the energeticsof impact ejecta emplacement were quite different. Thepioneering work summarized in Oberbeck et al. (1977) andOberbeck (1975) showed that in order for ejecta to betransported from the impact site to distances of tens to thousandsof kilometers, the ejecta must have a very high initial ejectionvelocity. Furthermore, due to the lack of an atmosphere,

deceleration was minimal, and the ejecta would reimpact withthe same velocity it left the initial cavity of excavation. Thefundamental implication of this fact was that ejecta would notbe emplaced as a passive "blanket", but that it would locallyimpact at very high velocity, excavating many multiples of itsown mass, to produce an ejecta deposit which would consistof a small component of the primary ejecta and a largecomponent of the locally excavated material. Thus, the final"ejecta deposit" would be emplaced in a very dynamic event,involving significant energetic erosion of the substrate, hugeclouds and surges of laterally moving ejecta, and relatedlandslides and debris flows, and ponding of ejecta. When the

FIG. 1. Feature near Discovery Scarp. This structure between thearrows has been suggested to be a dome, but identification is difficult.

Identification of mercurian volcanism 1211

event was over, the ejecta "deposit" would consist primarilyof locally excavated material, rather than an "ejecta blanket"of predominantly excavated material.

In a series of papers, Oberbeck and colleagues (seesummary in Oberbeck, 1975) showed that smooth plains couldreasonably be interpreted as accumulations of locally pondedejecta from such basin-scale impact events. In principle, thecomposition of these deposits should be characterized by amixture of primary and secondary ejecta material, and bedominated by secondary ejecta at increasingly greater distancesfrom the transient cavity due to the increasingly higher ejectavelocities at greater radial distances. Thus, ejecta spectralsignatures might be dominated by primary target material nearthe rim, but increasingly by local material excavated by thesecondary impacts and incorporated into the ejecta deposits atgreater radial ranges. Indeed, the detection of cryptomaria (Headand Wilson, 1992; mare volcanic deposits underlying higheralbedo ejecta deposits) showed the importance of mixing ofbasin ejecta with local material (e.g., Mustard and Head, 1996).

In most cases, there was evidence of ejecta emplacementby radial sculpting of highland terrain and smoothing ofintervening low terrain by ponded ejecta. Some very freshlunar craters, such as Tsiolkovsky, have lobate flow-likefeatures on their rims (e.g., Wilhelms, 1987, plate 3.24).However, there was little evidence of discrete flow fronts orlobes in association with the emplacement of these smoothplains at basin scales (e.g., Oberbeck, 1975), making difficultthe distinction between smooth plains formed by differentmechanisms such as effusive volcanism and ejecta ponding.In some cases there is evidence for lobate flow-like featuresthat appeared to represent late-stage movement andemplacement of these ejecta deposits. In the next section, weexamine the two most well-developed examples of these knownfrom lunar basins, and assess their origin, associated features,and the ability to recognize them under different conditions.

Lunar Examples from the Orientale Basin

One distinctive flow lobe has been identified ∼ 1000 kmsouth-southeast of the center of the Orientale Basin, in the midstof radial ejecta facies near the crater Inghirami (Moore et al.,1974, their Fig. 7; Wilhelms, 1987, his Fig. 4.4D). This flowlobe is ∼ 30 km wide, about 50–70 km long, and extends awayfrom the vicinity of the end of a major secondary crater chainradial to the Orientale Basin (Fig. 2). The flow lobe is ∼ 400 mthick and is convex outward away from the basin rim. It appearsto be composed of ejecta, rather than lava flows, and to havebeen formed in the terminal stages of the emplacement of theOrientale Basin ejecta deposit. The evidence for this includes(1) its crispness, (2) its superposition on ejecta plains and relateddeposits, (3) its spectral similarity to adjacent radially texturedejecta deposits, (4) its contrast to the spectral characteristicsof adjacent mare deposits, (5) its apparent topographic controlby newly formed ejecta topography, and (6) its scarp height

and thickness, much greater than any known lava flows on theMoon (e.g., Head and Wilson, 1992).

Clementine color composite images of the area show thatthe lobate flow is virtually invisible, being indistinguishablefrom the surrounding textured ejecta and smooth plains. LunarOrbiter and Clementine images (Fig. 2c,d) show that the originof the lobe blends with the more proximal ejecta facies andmay be related to the flow of mobilized ejecta from the end ofthe large secondary crater chain.

A second distinctive lobate flow is located just inside the craterGrimaldi (Fig. 3a–d). This lobate feature is similar to the onenear the crater Inghirami and has been described in Wilhelms(1987; his Fig. 4H). The structure itself is ∼ 35 km wide, up to∼ 60 km long, and is generally concave outward from the OrientaleBasin, with its local direction influenced by the presence ofGrimaldi. Its spectral characteristics are indistinguishable fromthose of the surrounding ejecta deposits and contrast distinctlywith those of the adjacent mare deposits on the floor of Grimaldi.This example too appears to have been derived from the late-stage movement of ejecta of the Orientale ejecta deposit.

Lessons for the Recognition of Ejecta-Related Plainson Mercury

On the basis of these examples and our general knowledgeof the emplacement of crater and basin ejecta (e.g., Oberbeck,1975), we can reach the following conclusions: (1) most ejecta-emplaced plains fill low-lying areas between radially textureduplands; (2) these plains can have multiple elevations due tothe elevation of the topography prior to the ejecta emplacementevent; (3) these plains distal to the basin will have mineralogiccharacteristics typical of a mixture of the projectile materialdominated by the local substrate; (4) in most cases, lobate ejectaboundaries are not seen in the smooth plains; (5) in some cases,ejecta lobes are seen, but they tend to be very localized and tobe related to late stage phases of the event; and (6) the spectralcharacteristics of these plains are virtually indistinguishablefrom the adjacent ejecta, but distinguishable from mare deposits.On the basis of these characteristics, we can further assess thenature of plains on Mercury and assess the possibility that theymay be of impact ejecta origin. Of course, in order to applythese observations effectively, one must keep in mind thedifferences in how the cratering processes and ejectaemplacement is manifested on Mercury due to variations inimpact velocities and gravity (e.g., Gault et al., 1977; Pike,1988; Schultz, 1988).

A comparison of Orientale Basin on the Moon with CalorisBasin on Mercury clearly demonstrates the morphologicaldifferences in ejecta emplacement on the two bodies. The flowfronts identified above are located in the portion of the Orientaleejecta blanket interpreted to be the Inner Facies of the HeveliusFormation. The Inner Facies contain elongate ridges andtroughs generally radial to the basin. Many of the ridges displaylobate edges similar to those of viscous flow fronts and are


oriented facing away from the central basin (McCauley, 1987).The Inner Facies are made up of the outward ground surge ofmaterial characteristic of an ejecta flow (McCauley, 1987).Surrounding the Inner Facies are the Outer Facies, made up ofrolling plains and weakly lineated terrain thought to be thedistal ends of the ejecta blanket. The Outer Facies arerecognizable up to one basin diameter out from Orientale(McCauley, 1977).

The outermost ejecta facies recognized around Caloris isthe Van Eyck Formation, a lineated terrain (Spudis and Guest,1988) proposed to be equivalent to the Inner Facies of Orientale(McCauley, 1977). The lineated terrain around Caloris isrecognizable out to a distance of one basin diameter (McCauley,1977), where it is embayed by smooth plains. These smoothplains are found radially outside of Caloris (Spudis and Guest,1988) and contain possible flow fronts (discussed below).

If the Van Eyck Formation and the Inner Facies are indeedequivalent, then ejecta emplacement on the Moon and Mercury

must be different. The Inner Facies do not extendproportionally as far from the associated basin as do the VanEyck Formation. This may imply that lobate flows travelfurther on Mercury than on the Moon. Alternatively, ejectaflows may have a different morphology on the two bodies dueto gravitational differences.

There are several reasons why the smooth plains encirclingCaloris Basin are unlikely to be crater ejecta deposits. Themost probable location for lobate ejecta deposits, based on a

FIG. 2. Ejecta flow lobe associated with Orientale Basin and near thecrater Inghirami. (a) Lunar Orbiter image of flow (subframe of imageIV-180-H1). (b) Sketch map of area in (a). (c) Clementine image offlow and surrounding deposits. (d) Lunar Orbiter image of flow andsurrounding deposits (image IV-180-H1).


comparison with Orientale Basin, is closer to the rim of Caloris(either in the Van Eyck Formation or other terrains closer tothe basin rim). The candidate flow fronts (discussed below)observed in the smooth plains are oriented in the wrongdirection—their scarps face towards the basin rim rather thanaway, as would be expected from an outward flow.

VOLCANISM

Theoretical Considerations of Mercurian Volcanism

We have discussed above how some have argued that theplains deposits could be basin ejecta, similar to those found atthe lunar Apollo 16 landing site (Wilhelms, 1976; Oberbeck etal., 1977), raising the possibility that there are no identifiable

volcanic units on Mercury (for a summary of the controversy,see Spudis and Guest, 1988). Fundamental questions remainabout the presence and nature of basic magmatic processes onMercury: Did Mercury, like the Moon (Head and Wilson, 1992;Head et al., 2000), form a primary crust which served as abarrier to, and filter for, magma ascent? Did secondary crustalformation (e.g., analogous to that of the lunar maria) occupy a

FIG. 3. Ejecta flow lobe associated with Orientale and located justinside the crater Grimaldi. (a) Lunar Orbiter image of flow (subframeof image IV-168-H3). (b) Sketch map of area in (a). (c) Clementineimage of flow and surrounding deposits. (d) Lunar Orbiter image offlow and surrounding deposits (image IV-168-H3).


significant part of the resurfacing history of Mercury? Did basalticvolcanism contribute to the resurfacing history of Mercury or didthe formation of the iron core so alter the mantle geochemistrythat other rock types dominate any eruptives? Does the tectonichistory of Mercury (which exhibits significant globalcompressional deformation) mean that extrusive volcanism wasinhibited or precluded by the state of stress in the lithosphere inits early to intermediate history? We now address these questionsthrough an assessment of the ascent and eruption of magma undermercurian conditions under a variety of settings.

It is commonly assumed that the large iron/silicate ratio ofMercury is the result of its early modification by a giant impact(Benz et al., 1988; Tyler et al., 1988; Tonks and Melosh, 1992).The extent of the differentiation of Mercury prior to this eventis moot as regards surface features and surface chemistry ifmost of the original crust and much of the original mantle werestripped away and lost into the Sun (Benz et al., 1988), andmuch of the remaining crust and mantle were left as a moltenmagma ocean (e.g., Benz et al., 1988; Tyler et al., 1988; Tonksand Melosh, 1992). However, the extent of any earlydifferentiation and the details of the impact event are significantfor the size of the iron core and for the bulk compositions ofthe residual unmelted mantle and the postulated overlyingmagma ocean. In particular, it is not automatically safe toassume that the mantle which is related to the present crustwas very close to chondritic in composition.

Whatever the composition of the mantle after the giantimpact, there is uncertainty about the extent to which solid-state convection could have proceeded in that mantle (Benz etal., 1988; Jeanloz and Morris, 1986; Spohn, 1991). The vigorof any convection that was occurring would have varied withtime as a function of the ever-decreasing amounts of long-lived radioactive heat sources and the ever-increasingtemperature difference across the mantle as a result of thecooling of the overlying magma ocean.

If convection in the mantle did take place, then partialmelting should have occurred and buoyancy forces should havebeen able to extract melts; these would have been of basalticcomposition if the mantle were chondritic but clearly othercompositions are possible (e.g., Jeanloz et al., 1995). Bodiesof buoyant melt would have risen, probably diapirically, untilthey encountered one or other of two kinds of trap (Head andWilson, 1992; Wilson and Head, 2001). The first option is adensity trap, that is, a level at which melts became neutrallybuoyant on encountering less dense overlying rocks: the baseof the crust is an obvious possibility for this. The second optionis a rheological trap, that is, a level at which the surroundingrocks are no longer able to deform in a plastic fashion fastenough to relax the stresses caused by the buoyancy forces.

Magma present in a rheological trap will still in general bebuoyant and may be able to rise further by opening a brittlecrack in the overlying rocks (i.e., by propagating a dike).However, any dike which forms will have a vertical extentlimited by the volume of magma available and the distribution

of stresses in the lithosphere (Head and Wilson, 1992; Wilsonand Head, 2001). Multiple such dikes may form (Head andWilson, 1992; Wilson and Head, 2001), but they will stall atshallower depths if they encounter density traps.

If the rheology of the rocks at all points around the peripheryof a magma body stalled at a rheological trap can supportstresses, an excess pressure may be created in the magma bodyby density changes consequent on chemical evolution. Suchevolution would be the consequence of fractional crystallizationwhich could conceivably lead to supersaturation of anyrelatively insoluble volatile phase and gas bubble formation,and this could aid dikes in penetrating much further into theoverlying rocks, thus overcoming density traps to some extent(Head and Wilson, 1992). A magma body stalled at a densitytrap may also acquire an excess pressure due to chemicalevolution, enabling it to migrate to shallower depths, if therocks surrounding it are elastic.

The optimum mechanism for allowing mantle melts denserthan the crust to penetrate to shallow depths as intrusions or toerupt at the surface is to have a rheological trap a short distancebelow a density trap (Wilson and Head, 2001). It is thenpossible for dikes to grow both upward into the crust anddownward into the mantle (which behaves elastically on theshort timescales of dike propagation (Wilson and Head, 2001)and to use the positive buoyancy of the melt in the mantle tooffset the negative buoyancy in the crust (Solomon, 1975).

How do these ideas relate to Mercury? If partial meltingbegan in the mantle while a magma ocean produced by a giantimpact was still liquid, any melts rising buoyantly through themantle that were also less dense than the magma ocean would,of course, have mingled with the ocean. Any such melts denserthan the ocean would have ponded beneath it, with perhapssome small amount of mechanical mixing occurring at theboundary. Only after the ocean had solidified could mantlemelts penetrate the resulting crust, and then only under thecircumstances described above.

The complete range of possible density structures forMercury after the solidification of any magma ocean, and theconsequent range of density contrasts between lithosphere andpossible mantle melts, is not very different from those that applyto the Moon. As a result the possible outcomes range from nomagmas reaching the surface, through magmas reaching thesurface only in topographic lows but otherwise stalling as(possibly shallow) intrusions, to circumstances in which largevolumes of magma reach the surface. The last option wouldbe favored by the fortuitous proximity of rheological and densitytraps near the crust–mantle boundary and could be related to thewidespread plains-forming units if these are in fact volcanic.Magmas formed by the chemical evolution of mantle melts inrheological traps would be the best candidates for the pyroclasticdeposits associated with the crater Homer. Ultimately these issueswill only be clarified by better identification of the compositionof surface units, currently rendered difficult both by the spatialand spectral resolution of available data (see following discussions)


and the complications due to space weathering of Mercury's surface(Hapke, 2001; Noble and Pieters, 2001).

Identifying Volcanism in Images

The Moon makes a useful analog for Mercury. Mercury'scrust is spectrally similar to the lunar highlands (Blewett et

al., 1997). Mare basalt deposits are the primary volcanic featureon the Moon and are often compared to the mercurian plains(e.g. Murray et al., 1975; Strom et al., 1975).

Lunar Volcanic Features–A wide variety of landformsare associated with the lunar maria (Head and Wilson, 1992).These include regional dark mantling deposits (Fig. 4a) whichare interpreted to be caused by pyroclastic eruptions. Glass

FIG. 4. Lunar volcanic landforms. (a) Dark mantling deposits on thefloor of Alphonsus (LO V-116-M). (b) Sinuous rilles. Hadley Rillefrom the Clementine Lunar Basemap. (c) Lava flow units. Largeflow in Mare Imbrium (LO V-160-H2,H1, V-161-H2,H1). (d) Shieldvolcanoes in Mare Insularum (LO IV-133-H1). (e) Volcanic complexin Marius Hills (LO V-211-M).


beads from these eruptions are spread tens to hundreds ofkilometers and are found in a number of lunar samples. Theeruptions must continue for a relatively long time to build upthe observed deposits. Thus the dark mantling deposits appearto mark locations of rapid and sustained eruptions (Head andWilson, 1991). Sinuous rilles (Fig. 4b) are often associatedwith these deposits as well as the maria. Sinuous rilles aremeandering channels generally an order of magnitude largerand often much more sinuous than terrestrial lava channelsand are found primarily in the lunar maria. They are thoughtto be sites of eruptions and thermal erosion. Sinuous rillesmay be the sites of the rapid and sustained eruptions whichproduced dark mantling deposits (Head and Wilson, 1991). Lavaflow units (Fig. 4c) are mapped on the basis of topographic,albedo, and color boundaries but individual flows are usuallyhard to identify due to impact degradation of individual flowunit boundaries (Head and Wilson, 1991). Flow fronts aremore easily identified in near-terminator images (Head andLloyd, 1973). About 50 small shield volcanoes (Fig. 4d) areidentified. They are characterized by a circular to irregularoutline and slopes <5°; nearly all of these features are found inassociation with mare. The shields range from 3 to 20 km inbasal diameter but no shields larger than ∼ 20 km are seen (Headand Gifford, 1980). Such volcanoes with diameters over 50 kmare common on Earth, Mars, and Venus; this difference isattributed to the lack of lunar shallow neutral buoyancy zonesaffecting eruption processes (Head and Wilson, 1991, 1992). Alsoconspicuous in their absence are calderas; these represent shallowmagma reservoirs at shallow neutral buoyancy zones. Finally,several volcanic complexes (Fig. 4e) with unusual concen-trations of the above features are seen; these are likely sites ofmultiple high-volume eruptions and the source regions of muchof the surrounding lava. The Marius Hills area (35 000 km2)contains 20 sinuous rilles and over 100 domes and cones whilethe Aristarchus Plateau/Rima Prinz region (40 000 km2)contains 36 sinuous rilles (Whitford-Stark and Head, 1977).We assess the detectability of such features underMESSENGER viewing geometry and lighting conditions.

Observations–Important volcanic sites were selected forthis study representing the range of volcanic features describedabove. The volcanic complexes of Marius Hills (∼ 12° N, 310° E)and Aristarchus Plateau/Rima Prinz (∼ 25° N, 315° E) containan assortment of features including sinuous rilles and volcanicdomes. The Apollo 15 landing site near Hadley Rille (∼ 25° N,3° E) is a classic example of a sinuous rille. Mare Imbriumcontains many lava flow fronts; the example used here is at∼ 33° N, 335° E. Additional volcanic shields are found in MareInsularum (∼ 7° N, 330° E). Images were taken from theClementine lunar basemaps at 100 and 500 m/pixel, and2.5 km/pixel resolutions. For comparison, the average resolutionof Mariner 10 images was 1.5 km/pixel while the highestresolution images were 90 m/pixel (Spudis and Guest, 1988).The anticipated average resolution for MESSENGER is250 m/pixel and the highest resolution images will be 25 m/pixel(Solomon et al., 2001).

Image Resolution–Figure 5 shows the three images for HadleyRille, a 120 km long, 1.5 km wide sinuous rille located at the edgeof Mare Imbrium. Figure 5a shows the rille at 100 m/pixel; herethe feature is clearly identifiable. At 500 m/pixel (Fig. 5b) thisrille is still easily identified, although its outline is pixelated.However, at 2.5 km/pixel resolution (Fig. 5c) the rille is notrecognizable. Resolutions of ~500 m/pixel (or slightly lower) arerequired to identify this feature.

Rima Prinz is a volcanic complex containing an extensiveseries of sinuous rilles (Fig. 6). Four of these rilles are markedwith arrows in Fig. 6a; the rilles are clearly recognizable in the100 m/pixel image. In the 500 m/pixel image (Fig. 6b), two ofthese rilles are still identifiable; however the smaller rilles(upper left and furthest right) are much harder to locate. Finally,at 2.5 km/pixel none of the rilles can be recognized. Resolutionsbetween 100 and 500 m/pixel are needed for rille identification.

Figure 7 contains the same volcanic domes in MareInsularum as Fig. 4d; they are much more difficult to identifyin the Clementine lunar basemap image due to their differentSun elevation angles (see following section on viewinggeometry). Arrows in Fig. 7a point to the locations of these

FIG. 5. Hadley Rille (∼ 25° N, 3° E) at multiple resolutions. (a) 100 m/pixels. (b) 500 m/pixels. (c) 2.5 km/pixels. All images from theClementine lunar basemap.


domes; the leftmost dome is easiest to identify based on albedodifferences with the surrounding material. This dome isrecognizable in the 100 m/pixel image (Fig. 7a). In the 500 m/pixel image (Fig. 7b) it is difficult to distinguish the dome fromother albedo variations in the area, and in the 2.5 km/pixelimage none of the features in this region are identifiable. Toidentify small domes, resolutions between 100 and 500 m/pixelare necessary.

Flow fronts are also very difficult to identify in theClementine lunar basemap. Figure 8 contains the sameImbrium flow front found in Fig. 4d. The flow front is outlinedin Fig. 8a based on close comparison with low Sun LunarOrbiter images. In several locations the edge of the flowcorresponds to albedo boundaries. While the flow front itselfis not visible in any resolution figure, these boundaries arerecognizable in both the 100 and 500 m/pixel images (Fig. 8b

FIG. 6. Rima Prinz at multiple resolutions. (a) 100 m/pixels. (b) 500 m/pixels. (c) 2.5 km/pixels. All images from the Clementine lunarbasemap.

FIG. 7. Domes located in Mare Insularum at multiple resolutions. (a) 100 m/pixels. (b) 500 m/pixels. (c) 2.5 km/pixels. All images from theClementine lunar basemap.

FIG. 8. Flow front in Mare Imbrium at multiple resolutions. (a) 100 m/pixels. (b) 500 m/pixels. (c) 2.5 km/pixels. All images from theClementine lunar basemap.


and 8c, respectively). However, the 2.5 km/pixel image is sopixelated that the boundaries are unreliable.

Viewing Geometry–As mentioned in the above discussionof the Imbrium flow front and Insularum domes, viewinggeometry plays a role in identification of volcanic features.This is also easily seen at Marius Hills, a volcanic complexwith numerous, closely spaced domes as well as sinuous rilles,rimless depressions thought to be of volcanic origin, and mareridges (Guest, 1971). Figure 9 shows two ground-basedtelescopic images of Marius Hills taken at different co-longitudes. The image in Fig. 9a has a higher Sun elevationangle (26°; Malin, 1978) and shows faint topographic features;however, it is difficult to interpret these features. In comparison,Fig. 9b was taken at low Sun elevation (6°; Malin, 1978) andclearly contains domes and ridges. The only difference betweenthese images is Sun elevation. Thus, lighting conditions arean important factor in identifying volcanic features in images.

To summarize, in order to identify a volcanic feature on aplanetary body, it is necessary to have 2–3 pixels containingthat feature; thus the necessary resolution for identification isa function of landform size. For typical lunar volcaniclandforms, resolutions of 100–500 m/pixel are observed to benecessary for identification. Larger features such as sinuousrilles can be observed at 500 m/pixel, while smaller featuressuch as domes require higher resolutions closer to 100 m/pixel.Additionally, many features such as domes and flow frontsrequire low Sun elevation angles in order to be identified. Inorder to unequivocally conclude that volcanism has beenpresent in the history of a planet several clearly volcanicfeatures ought to be identified.

CANDIDATE FLOW FRONTS ON MERCURY

Although the majority of Mariner 10 images are at resolutionsand Sun elevation angles inappropriate for identifying volcanism,a small number of images were taken at <500 m/pixel and near

the terminator (low Sun angle). An examination of these imagesreveal possible flow fronts which can be interpreted to have formeda number of ways. This demonstrates the difficulty of definitivelyidentifying volcanic features.

Lobate scarps are relatively steep escarpments tens to hundredsof kilometers in length found across the surface of Mercury. Theyare characterized by rounded crests and lobate outlines on thescale of a few to tens of kilometers. 17% of observed scarpsappear on the smooth plains on crater floors and form boundariesbetween different terrains (Strom et al., 1975); Dzurisin (1978)classified these scarps as irregular intracrater scarps. These scarpshave been interpreted as flow fronts (Strom et al., 1976) and astectonic features (Dzurisin, 1978). We examine three examplesof candidate lobate flow fronts in Mariner 10 data and discussvarious options for interpretation. These three examples comefrom the Shakespeare Quadrangle of Mercury (H-3) as mappedgeologically by J. E. Guest and R. Greeley (1983) and are arrayedto the east of the Caloris Basin rim.

The first example (Fig. 10) is a 180 m/pixel image and themost distal from the Caloris Basin rim (∼ 600 km). It showsthe contact of smooth plains (right) embaying the Van EyckCrater rim (left), composed of radially textured Odin Formation,a facies of the Caloris Basin ejecta material. The contact(location denoted by arrows) is sharp, and the smooth plainsappear to embay the inner base of the crater rim. The westernmargin of the smooth plains locally consists of a westward-facing scarp that is lobate and flow-like and appears to conformto preexisting topography. This flow could be interpreted aseither volcanic or as an ejecta flow. However, this may not bea flow at all; in some places it is also similar to wrinkle ridgesin morphology. If the entire feature were made up of wrinkleridges then this would be a tectonic feature formed bycompressive stresses in the region. This last possibility seemsunlikely due to the embayment relations between the flow andthe preexisting topography; a wrinkle ridge is more likely tocut through preexisting topography.

FIG. 9. Two views of Marius Hills (~12° N, ~310° E) from ground-based telescopes. From The Consolidated Lunar Atlas (Kuiper et al.,1967). (a) Image C790. Sun angle is 26°. (b) Image C4268. Sun angle is 6°.


The second example (Fig. 11) is a 280 m/pixel image and is atan intermediate distance from the Caloris Basin rim (∼ 450 km).It shows the contact of smooth plains (lower middle of image)with another subunit of the Odin Formation (hummocky plains),a facies of the Caloris Basin ejecta material (left and upper part ofimage). The western margin of the smooth plains (denoted bywhite arrows) appears to embay the Odin Formation and to floodand cover the hummocks that characterize it in this region. Thecontact consists of a westward-facing scarp that is sinuous in natureand appears to embay local topography. This is a particularlygood example to illustrate both: (1) the lobate, flow-like natureof the contact, and (2) its morphological similarity to wrinkle ridgesthat are pervasive in the smooth plains (lower center and right inimage). Could this contact be a tectonic wrinkle ridge? Again, ifthis feature is a flow front it is impossible to tell if it is a lava flowfront or a crater ejecta flow front.

The third example (Fig. 12) is a 300 m/pixel image and isclosest to the Caloris Basin rim (∼ 230 km). This image showstwo westward-facing lobate scarps (white arrows) in the NervoFormation (hummocky, rolling plains), a facies of the CalorisBasin ejecta interpreted to be fallback mixed with impact melt.The surrounding hills consist of the Odin Formation, agradational facies of Caloris Basin ejecta. Patches of smoothplains are seen in the upper right and lower middle. Themorphology of the lobate scarps is strongly suggestive of flowand embayment, but also bears similarities to wrinkle ridgemorphology (compare to Fig. 10). It is also difficult todistinguish the morphology from lava flow fronts in theImbrium Basin on the Moon (Fig. 2c) and Orientale Basin ejectaflow lobes also on the Moon.

Through this examination of candidate flow front featureson Mercury, we see that in all cases it is impossible to tell if

FIG. 10. Smooth plains (center of image) embaying the inner rim of crater Van Eyck. Arrows indicate flow-like scarps. Mariner 10 image0000077, 180 m/pixel.


these features are due to volcanic or crater ejecta flows.Compositional data is needed to distinguish between types offlows; if the flow front also marks the boundary of two distinctcompositional regions, then the flow is volcanic. Identificationcannot be based on morphology alone.

IMPLICATIONS FOR MESSENGER

One of the major scientific goals of the MESSENGERmission is to determine the formation history of the surfaceof Mercury (Solomon et al., 2001). An important aspect ofthis is the role of volcanism in Mercury's history;MESSENGER needs to be able to search for volcaniclandforms under suitable resolutions and viewing geometriesto correctly address this issue. This requires adequatecoverage of the planet at the resolutions and viewinggeometries outlined above.

The mission timeline and the anticipated image resolutionscan be found in Solomon et al. (2001). The spacecraft camerasystem will build up a full global mosaic in the first 6 monthsafter achieving orbit and cover the planet with stereomonochrome images at ∼ 250 m/pixel in the second 6 months.

Two flybys prior to orbit insertion will allow 85% of the planetto be imaged in monochrome at 500 m/pixel and in color at anaverage resolution of ∼ 2.4 km/pixel. Since the spacecraft willbe in an elliptical orbit, the narrow angle imager will be usedduring high altitude to achieve global monochrome imagemosaics with an average 250 m/pixel resolution and a highestresolution of 25 m/pixel. Color images will average 1.1 km/pixelwith targeted high-resolution color images at ∼ 300 m/pixel. Theorbit itself starts in a dawn–dusk configuration and movesthrough noon–midnight back to dawn–dusk many times duringthe mission (Solomon et al., 2001).

Will this allow volcanic landforms to be identified? Mostof the imaged 45% of the planet is covered with plains thatmay be volcanic in origin (Spudis and Guest, 1988). In thefirst 6 months the spacecraft will cover ∼ 50 % of the surfacefrom near dawn–dusk orbits. Even if the only plains of possiblevolcanic origin are the ones already observed, this orbitgeometry will allow imaging of a portion of such plains at∼ 250 m/pixel and with low Sun elevation angles. These areconditions favorable to the identification of volcanic landforms.Additionally, repeat coverage of the planet will allow fortargeting of possible volcanic features with the high-resolution

FIG. 11. Smooth plains (lower middle of image) in the Odin Formation. Arrows indicate margin of plains which appears to embay the OdinFormation. Mariner 10 image 0000104, 280 m/pixel.


camera. Furthermore, MESSENGER multispectral imagingand spectrometer data will permit further assessment andcharacterization of candidate volcanic plains (Solomon et al.,2001; Robinson and Taylor, 2001) as well as distinguishbetween possible volcanic flow fronts and lobate ejecta flows.If volcanism exists on Mercury's surface, MESSENGER iswell-equipped to find it.

CONCLUSIONS

The formation mechanism of the plains of Mercury hasbeen debated since Mariner 10 revealed their existence. Bothcrater ejecta flows and volcanism have been suggested asmechanisms. Definitive identification of volcanism on Mercuryrequires concrete evidence such as an image of an unequivocalvolcanic landform. No such image has been found; however,this does not mean that volcanism has been ruled out. Bothresolution and viewing geometry play an important role in whatsorts of features are visible in an image. To successfully identifyvolcanic features such as those on the Moon both low Sunangle (∼ 10°) and ∼ 500 m/pixel resolution or better are needed.

While Mariner 10 did not cover a significant portion ofMercury within these requirements, MESSENGER will be ableto do so. The question of the existence of volcanism on Mercurymust wait until then to be addressed more confidently.

Acknowledgements–Thanks are extended to Barbara Cohen andLaszlo Keszthelyi for helpful reviews which improved the manuscript.We gratefully acknowledge support from the NASA PlanetaryGeology and Geophysics program and the MESSENGER mission toJ. W. H.

Editorial handling: W. K. Hartmann

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