E-type asteroid spectroscopy and compositional modeling - CiteSeerX

JOURNAL OF GEOPHYSICAL RESEARCH, VOL. 109, E02001, doi:10.1029/2003JE002200, 2004

E-type asteroid spectroscopy and compositional modeling Beth Ellen Clark,1 Schelte J. Bus,2 Andrew S. Rivkin,3 Timothy McConnochie,4 Josh Sanders,1 Sweta Shah,1 Takahiro Hiroi,5 and Michael Shepard6 Received 20 October 2003; revised 28 November 2003; accepted 5 December 2003; published 5 February 2004.

[1] We present near-infrared spectroscopic observations (0.8–2.5 mm) of E-type asteroids.

We combine these observations with visible wavelength spectra obtained by other researchers and perform Hapke theory mixing model simulations of E-type asteroid spectra in order to constrain possible compositions. Aubrites were originally suggested as the meteorite analog for the E-type asteroids because of their similar visible wavelength colors and high albedos. The designation ‘‘E’’ was originally linked with the mineral enstatite, common in aubrite meteorites. More recently, the sulfides troilite and oldhamite have been suggested as possible components of E-type asteroids. We tested the suggested compositional interpretations of aubrite meteorites and/or aubrites enhanced with sulfides. We also tested compositions of aubrites mixed with low-iron silicate minerals. We find that E types can be separated into three groups on the basis of inferred composition: ‘‘Nysa-like’’ E types are consistent with silicate mineralogy higher in iron than the mineral enstatite; ‘‘Angelina-like’’ asteroids are consistent with silicate mineralogy, including a sulfide such as oldhamite; and ‘‘Hungaria-like’’ E types are not inconsistent with aubrites. Our results indicate that some E-type asteroids may be composed of materials that are not INDEX TERMS: 6205 Planetology: Solar System Objects: Asteroids and sampled by meteorites. meteoroids; 6240 Planetology: Solar System Objects: Meteorites and tektites; 5410 Planetology: Solid Surface Planets: Composition; 6060 Planetology: Comets and Small Bodies: Radiation and spectra; 6055 Planetology: Comets and Small Bodies: Surfaces and interiors; KEYWORDS: asteroids, meteorites, asteroid-meteorite connection, reflectance spectroscopy, compositional modeling Citation: Clark, B. E., S. J. Bus, A. S. Rivkin, T. McConnochie, J. Sanders, S. Shah, T. Hiroi, and M. Shepard (2004), E-type asteroid spectroscopy and compositional modeling, J. Geophys. Res., 109, E02001, doi:10.1029/2003JE002200.

1. Introduction [2] Currently, a state of confusion exists as to the composition and meteorite linkage of the E-type asteroids. E types were originally described by Zellner [1975], Zellner and Gradie [1976], and Zellner et al. [1977] on the basis of observations of the very high albedo objects 44 Nysa and 64 Angelina (see Table 1). Using these designations, Tholen [1984] and Zellner and Tholen [1985] showed that approximately 14 of the 589 objects that they surveyed possessed similar photometric colors (0.3 –1.0 mm) and high albedos. More recently, the E-type class has grown to include 1

Department of Physics, Ithaca College, Ithaca, New York, USA. Institute for Astronomy, University of Hawaii at Hilo, Hilo, Hawaii, USA. 3 Department of Earth, Atmospheric and Planetary Sciences, Massachusetts Institute of Technology, Cambridge, Massachusetts, USA. 4 Department of Astronomy, Cornell University, Ithaca, New York, USA. 5 Department of Geological Sciences, Brown University, Providence, Rhode Island, USA. 6 Department of Geography and Earth Sciences, Bloomsburg University of Pennsylvania, Bloomsburg, Pennsylvania, USA. 2

Copyright 2004 by the American Geophysical Union. 0148-0227/04/2003JE002200$09.00

23 objects [Tholen, 1989; Gaffey et al., 1992; Gil-Hutton and Benavidez, 2001; Binzel et al., 2002; Kiselev et al., 2002; Delbo et al., 2003]. Table 1 lists all E types found up to this point in time. [3] As seen in Table 1, the Tholen E-class asteroids correspond roughly but imperfectly to the Xe class in the Bus taxonomy [Bus, 1999; Bus and Binzel, 2002b]. There are almost as many asteroids which are E but not Xe (and vice versa) as there are asteroids which are both Xe and E. Since the two taxonomies use slightly different criteria and different data sets (Tholen used albedo; Bus did not), it is not surprising that their members do not all correspond one to one, particularly for a relatively rare class like the E class. However, the rough connection between the E and Xe classes may lead to conclusions about both classes which are only valid for one. [4] It is evident from Table 1 that E types can perhaps be separated dynamically into two main groups: those in the Hungaria region from 1.88 to 1.98 AU and those beyond, from 2.10 to 2.72 AU. Tholen [1984] discusses the possible significance of family membership in the Hungaria region. Other researchers have extensively studied families in this region [Bendjoya and Zappala, 2002; Valsecchi et al., 1989; Zappala et al., 1995], but the consensus seems to be that the Hungaria asteroids are not a dynamical family.

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Table 1. Known E-Type Asteroids Asteroid

Name

44 64 214 317 434 504 620 1025 1103 1251 2035 2048 2449 2577 3101 3103 4660 5751 5806 6435 6911 7579 33342

Nysa Angelina Aschera Roxane Hungaria Corac Drakonia Riema Sequoia Hedera Stearns Dwornik Kenos Litva Goldberger Eger Nereusd Zaod Archieroy Daveross Nancygreen 1990 TN1 1998 WT24e

Semimajor Axis, AU IRAS Albedo Geometric Albedoa Polarization Albedoa Tholen Class Bus Class Bands Detected,b mm 2.424 2.683 2.611 2.287 1.944 2.725 2.435 1.979 1.934 2.717 1.884 1.954 1.901 1.900 1.979 1.405 1.989 2.104 1.963 1.919 1.932 1.978 0.718

0.49 0.43 0.52

0.55

0.41 0.48

0.52 0.49

0.46 0.34

0.48

E E E E E E E E E E E E E E

0.48 0.41

Xe Xe Xc Xe Xe X

[0.5], 0.9, 1.8, 3.0 0.5, 0.9, [1.8], 3.0 [0.5], 0.9, 1.8, 3.0 [0.5], 0.9, 1.8, 3.0 0.5, 0.9, [1.8] [3.0]

Xe Xk X Xe

[0.5], 0.9, [1.8] [0.5], 0.9, [1.8] 0.4, 0.5, [3.0] 0.9, [1.8]

0.5, 0.9

0.43

a

Geometric and polarization albedo values are from Tedesco et al. [2002] and Shevchenko et al. [2003], respectively. b Brackets mean that data exist at this wavelength but a band was not detected. c E-type designation for this asteroid is suggested on the basis of the X classification by Bus and Binzel [2002b] and the high albedo from Tedesco et al. [2002]. d These E-type designations were recently suggested by Delbo et al. [2003]. Asteroid 3103 Eger is suggested to be an E type by Gaffey et al. [1992]. Asteroids 5806 through 7579 were suggested to be E types by Gil-Hutton and Benavidez [2001]. Tholen E-class objects for which no albedos are given were designated on the basis of unpublished albedos provided by J. Gradie and E. F. Tedesco (D. J. Tholen, personal communication, 1984). e Polarization albedo and suggested E-type designation are from Kiselev et al. [2002].

[5] E types were originally proposed to be linked to the enstatite achondrite meteorites (also known as the aubrites) by Zellner et al. [1977], a link which was later explored in particular for a near-Earth asteroid, 3103 Eger [Gaffey et al., 1992]. In fact, in most summaries of inferred meteoriteasteroid links, one finds that this link between E-class asteroids and aubrite meteorites is well established in the literature [e.g., Bell et al., 1989; Gaffey et al., 1989; Cloutis et al., 1990, Cloutis and Gaffey, 1991, 1993; Pieters and McFadden, 1994; Fornasier and Lazzarin, 2001; Burbine et al., 2002a; Kelley and Gaffey, 2002]. [6] Recent visible wavelength spectra of E types have been published by Bus and Binzel [2002a, 2002b] and Fornasier and Lazzarin [2001]. (Kelley and Gaffey [2002] present a visible wavelength spectrum of 434 Hungaria only.) It was Bus who first noted a very unusual absorption feature from 0.44 to 0.51 mm in spectra of asteroids 64 Angelina, 434 Hungaria, and 3103 Eger [Burbine et al., 1998]. The feature (which we will call a 0.5-mm band) had been missed by the broadband passes used by Zellner and Tholen [1985] but was later confirmed by Fornasier and Lazzarin [2001], who also found the feature to exist in the spectrum of 2035 Stearns (and to be absent in spectra of 317 Roxane and 1251 Hedera). Kelley and Gaffey [2002] dispute the reality of the feature and propose that it may be related to UV fluorescence of certain enstatite crystals. Burbine et al. [1998] proposed that the feature is due to troilite, but this is problematic because E types are by definition very high albedo objects and troilite is naturally a very dark mineral component. Clark and Lucey [1984] showed that even minute amounts of a spectrally dark component can nonlinearly darken a mineral mixture, making it very unlikely that troilite is a major component of the surfaces of these bright objects. Nevertheless, troilite is a

component found in minor amounts in aubrite meteorites [Watters and Prinz, 1979; Keil, 1989; Burbine et al., 1998]. Finally, Burbine et al. [2002b] propose that the 0.5-mm band may be due to oldhamite, a mineral found in aubrites. [7] The visible wavelength spectra of E types are degenerate with (statistically not distinguishable from) M-, P-, and EMP- (or X) class objects [Tholen, 1984]. In the 52-color survey data of Bell et al. [1988], there are low signal-tonoise spectra of E-type asteroids 44 Nysa, 64 Angelina, and 317 Roxane. [8] Rivkin et al. [1995] and Rivkin [1997] find that four out of six E types surveyed show possible water-of-hydration absorption features near 3 mm (see Table 1). The two E types observed that did not show 3-mm bands were 620 Drakonia and 2035 Stearns. Such 3-mm features were not expected on E types on the basis of the presumed connection with the aubrite meteorites. Hydrated mineralogy could not survive the heating necessary to form the igneous rocks of the aubrite meteorites. Alternative explanations for E-type composition are summarized by Fornasier and Lazzarin [2001], but none of the proposed compositions is completely consistent with a known meteorite type. In particular, the 0.5-mm absorption band found in 64 Angelina, 434 Hungaria, 3103 Eger, and 2035 Stearns does not exist in aubrite meteorites. [9] The data obtained thus far that are relevant to constraints on the composition of the E-type asteroids are the absorption features and albedos summarized in Table 1. Table 1 updates a similar table published by Fornasier and Lazzarin [2001] and indicates that E types have spectral absorptions at several different wavelengths, none of which is found in aubrite meteorites. It appears unlikely that any particular aubrite can be linked with E-type asteroids, but could a mixture of aubrites possibly simulate E types more

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Figure 1. Spectra of asteroids 44 Nysa, 64 Angelina, 214 Aschera, 317 Roxane, 434 Hungaria, 1103 Sequoia, 1251 Hedera, 2048 Dwornik, and 3103 Eger. The visible wavelength region observations of 2048 Dwornik were obtained by Zellner and Tholen [1985]. All other visible spectra were obtained by Bus and Binzel [2002b]. Infrared observations were obtained at the NASA Infrared Telescope Facility on Mauna Kea; the spectrum of 3103 Eger was provided by T. H. Burbine and R. P. Binzel (T. H. Burbine, 2000, personal communication). (left) Combined spectra, scaled to 1.0 at 0.55 mm and offset from each other for clarity. (right) Infrared spectra only, after division by a straight-line fit to the spectral continuum. Note the 0.9-mm absorption features in 44 Nysa, 214 Aschera, and 317 Roxane. The same asteroids show strong evidence of a 1.8-mm band. These features are attributable to iron-bearing pyroxene silicate minerals. effectively? In this paper, we present new spectral measurements of nine E-type asteroids and use a Hapke theory compositional mixing model to simulate aubrite mixtures and to constrain possible meteorite and/or mineral interpretations of the range of E-type reflectance spectra.

2. Observations [10] The data presented here are combinations of visible wavelength observations [Bus and Binzel, 2002a, 2002b;

Zellner and Tholen, 1985] and infrared wavelength observations that we obtained (see Figure 1). For details on the visible wavelength observing circumstances, please see Bus and Binzel [2002a] and Zellner and Tholen [1985]. The infrared observations were obtained at Mauna Kea Observatory in Hawaii with the 3.0-m NASA Infrared Telescope Facility equipped with a cooled grating and an InSb array (1024
1024) spectrograph (SpeX) [see Rayner et al., 2003]. [11] Spectra were recorded with a slit oriented in the eastwest direction and opened to 0.8 arc sec. A filter cutting the

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Table 2. E-Type Asteroids Observed for This Study Asteroid

Name

Semimajor Axis, AU

Albedoa

mv

Date of Observation

Solar Analog

44 64 214 317 434 1103 1251 2048 3103

Nysa Angelina Aschera Roxane Hungaria Sequoia Hedera Dwornik Egerc

2.424 2.683 2.287 2.287 1.944 1.934 2.717 1.954 1.405

0.48 0.46 0.52 0.49 0.47 0.48 0.41 0.46b 0.46b

12.3 11.7 14.2 14.1 13.4 14.1 14.8 15.6

2001 2001 16 Aug. 2003 14 March 2001 2001 16 Aug. 2003 18 Aug. 2003 16 Aug. 2003

Landolt Landolt Landolt 107 – 684 SAO 65083 Landolt Landolt 112 – 1333 SAO 65083 Landolt 112 – 1333

a Albedos presented here are averages of the published values listed in Table 1. These are the values we used for the mixing model study. b Assumed value was calculated from the average of published E-type albedos. c The 3103 Eger spectrum is contributed by R. P. Binzel and T. H. Burbine (personal communication, 2003). See Burbine [2000] for details of the observations.

signal below 0.8 mm was used to prevent order overlap. Table 2 lists the asteroids and observing circumstances for the objects observed for this study. [12] Following normal data reduction procedures of flat fielding, sky subtraction, spectrum extraction, and wavelength calibration, each spectrum was fit with the ATRAN model for telluric absorption features [Lord, 1992]. This procedure required an initial estimate of precipitable water in the atmospheric optical path. We calculated this using the zenith angle for the observation and the known tau values (average atmospheric water) for Mauna Kea Observatory in Hawaii. This initial guess was iterated until the best fit between predicted and observed telluric band shapes was obtained and an atmospheric model spectrum was generated. Next, each asteroid spectrum was divided by the atmospheric model and was ratioed to each star spectrum (similarly reduced) before normalization at 1.0 mm. Solar analog star spectra were used, as listed in Table 2, and each star was measured three to five times per night. Thus we usually obtained three to five asteroid/ star ratios for each object, minimizing chances of spurious measurements. The final spectra that we report are averages of all ratios obtained for each object. Error bars are not plotted, but formal errors propagated through the reduction are less than the scatter in the data, so we assume that the scatter in the data is the best estimate of the uncertainties of each spectrum measurement.

3. Compositional Model [13] To set constraints on the possible composition of the E-type asteroids, we have constructed a mixing model that is based on Hapke’s theory of the reflectance of particulate surfaces [Hapke, 1981, 1984, 1986, 1993]. Our treatment of the Hapke model follows that of Mustard and Pieters [1989] and of Clark [1995] and references therein. Assuming reasonable end-member mineral choices, we use our mixing model to simulate the conditions of intimately mixed particles at the surface of an asteroid and then to minimize the difference between a model mixture and the asteroid spectrum. In particular, we wish to test the hypothesis that a mixture of aubrites (or aubrites and sulfides) can simulate the spectra of E-type asteroids. [14] The surface of an atmosphereless planet exposed to the impacts of interplanetary meteor bombardment is generally thought to be composed of fine-grain particles that are

mixtures of the substrate rock minerals. Although the particle sizes on asteroid surfaces are not known, evidence suggests that grain sizes will be coarser than on the Moon. The optically dominant lunar regolith particles range in grain size from about 10 to 50 mm, so it is reasonable to assume that grain sizes on asteroids will vary between 30 and 300 mm. Because finer-grain particles tend to coat coarser-grain particles, rendering the finer grains of the mixture more optically important, we decided to use grain sizes at the lower end of the assumed asteroid grain-size range, and we set our grain sizes at less than 125 mm. [15] Neglecting factors that account for macroscopic surface roughness, the first equation of our Hapke model is
  w bl cosðaÞ þ cl 1:5 cos2 ðaÞ  0:5 4ðmo þ mÞ    1 þ 2m 1 þ 2mo pffiffiffiffiffiffiffiffiffiffiffiffiffiffi pffiffiffiffiffiffiffiffiffiffiffiffiffiffi ; þ 1 þ 2mo 1  wl 1 þ 2m 1  wl

rl ¼

where a denotes the effective phase angle [Hapke, 1981, equation (37)] and w denotes single-scatter albedo. The parameters m = cos(emission angle) and mo = Cos(incidence angle) are known from the relative positions of the observer, target, and illumination source at the time of the observation. The result, rl, is the bidirectional radiance coefficient, defined as the ratio of the bidirectional reflectance of a surface to that of a perfectly diffuse surface illuminated at the same angle of illumination as the sample [Hapke, 1981, 1993]. [16] The parameter wl is the single-scatter albedo of an individual particle. The single-scatter albedo is the energy scattered by a particle at a given wavelength as a fraction of the incident energy at that wavelength in a single-particlescattering event. Hapke theory says that the single-scatter albedo for a mixture is the average of the single-scatter albedos of the mixture components, weighted by the abundances of the components. The H functions are the Chandrasekhar scattering functions for isotropic particles. These functions are described by Hapke [1981]. The physical meaning of Hapke parameters is more fully discussed by Hapke [1981, 1984, 1986, 1993], Helfenstein and Veverka [1989], and Helfenstein et al. [1994, 1996]. [17] This form of the Hapke function uses an empirical description of the single-particle scattering phase function, a Legendre polynomial, with coefficients cl and bl measured

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Table 3. Mineral and Meteorite End-Members RELAB File Name

Name

Mix Code

Grain Size, mm

Type

Source

c1pp21 c2pd01 c1po76 c1pe30 c1tb38 cltb55 cdmb06 c1mt40 c1tb47 cbea03 s1tb46 cbea04 c1tb45

clinopyroxene diopside forsterite orthopyroxene oldhamite sulfide troilite Abee Bishopville Happy Canyon Mayo Belwa Norton County Pena Blanca

CPX DIO FOR OPX OLD SUL TRO ABE BIS HAP MAY NOR PEN