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Icarus 170 (2004) 259–294www.elsevier.com/locate/icaru

Observed spectral properties of near-Earth objects: results for popudistribution, source regions, and space weathering processes

Richard P. Binzela,∗, Andrew S. Rivkina, J. Scott Stuarta, Alan W. Harrisb, Schelte J. Busc,Thomas H. Burbined

a Department of Earth, Atmospheric, and Planetary Sciences, Massachusetts Institute of Technology Cambridge, MA 02139, USAb Space Science Institute, 4603 Orange Knoll, La Canada, CA 91011, USAc Institute for Astronomy, 640 North A’ohoku Place, Hilo, HI 96720, USA

d Laboratory for Extraterrestrial Physics, NASAGoddard Space Flight Center, Greenbelt, MD 20771, USA

Received 20 November 2003; revised 20 March 2004

Available online 11 June 2004

Abstract

We present new visible and near-infrared spectroscopic measurements for 252 near-Earth (NEO) and Mars-crossing (MC) objecfrom 1994 through 2002 as a complement to the Small Main-Belt Asteroid Spectroscopic Survey (SMASS,http://smass.mit.edu/). Combinedwith previously published SMASS results, we have an internally consistent data set of more than 400 of these objects for investigating trrelated to size, orbits, and dynamical history. These data also provide the basis for producing a bias-corrected estimate for thepopulation (Stuart and Binzel, 2004, Icarus 170, 295–311). We find 25 of the 26 Bus (1999, PhD thesis) taxonomic types are rewith nearly 90% of the objects falling within the broad S-, Q-, X-, and C-complexes. Rare A- and E-types are more common in theNEO population (about 5% compared to< 1%) and may be direct evidence of slow diffusion into MC orbits from the Flora and Hunregions, respectively. A possible family of MC objects (C-types) may reside at the edge of the 5:2 jovian resonance. Distinct signrevealed for the relative contributions of different taxonomic types to the NEO population through different source regions. E-typan origin signature from the inner belt, C-types from the mid to outer belt, and P-types from the outer belt. S- and Q-types have eidentical main-belt source region profiles, as would be expected if they have related origins. A lack of V-types among Mars-crosserentry into NEO space via rapid transport through theν6 and 3:1 resonances from low eccentricity main-belt orbits, consistent with aorigin. D-types show the strongest signature from Jupiter family comets (JFC), with a strong JFC component also seen among thA distinct taxonomic difference is found with respect to the jovian Tisserand parameterT , where C-, D-, and X-type (most likely lowalbedo P-class) objects predominate forT � 3. These objects, which may be extinct comets, comprise 4% of our observed sample, blow albedos makes this magnitude limited fraction under-representative of the true value. With our taxonomy statistics providingcomponent to the diameter limited bias correction analysis of Stuart (2003, PhD thesis), we estimate 10–18% of the NEO populaany given diameter may be extinct comets, taking into account asteroids scattered intoT < 3 orbits and comets scattered intoT > 3 orbits.In terms of possible space weathering effects, we see a size-dependent transition from ordinary chondrite-like (Q-type) objects to S-typeasteroids over the size range of 0.1 to 5 km, where the transition is effectively complete at 5 km. A match between the average sof 5 km asteroids and the rate of space weathering could constrain models for both processes. However, space weathering maya very rapid rate compared with collisional timescales. In this case, the presence or absence of a regolith may be the determininwhether or not an object appears “space weathered.” Thus 0.1 to 5 km appears to be a critical size range for understanding thetimescales, and conditions under which a regolith conducive to space weathering is generated,retained, and refreshed. 2004 Elsevier Inc. All rights reserved.

Keywords:Asteroids; Asteroids composition; Surfaces asteroids; Asteroids near-Earth

r inEO)-

* Corresponding author. Fax: 617-253-2886.E-mail address:[email protected] (R.P. Binzel).

0019-1035/$ – see front matter 2004 Elsevier Inc. All rights reserved.doi:10.1016/j.icarus.2004.04.004

1. Introduction

Many pieces of the puzzle must be brought togetheorder to have a clear picture of the near-Earth object (Npopulation. Four of the piecesthat can be described in

http://www.elsevier.com/locate/icarus

http://smass.mit.edu/

260 R.P. Binzel et al. / Icarus 170 (2004) 259–294

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clude: (i) the taxonomic distribution of the populationmeasured by observational sampling, (ii) the determinaof albedos that can be associated with the taxonomictribution, (iii) discovery statistics for the NEO populatioand (iv) the debiasing of the discovery statistics usingtaxonomic and albedo information. This paper presentsfirst piece, detailing the observations and observed chateristics of the NEO and Mars-crossing (MC) populatiFor the second piece, a complementary program of albmeasurements was pursued at the Keck Observatory (BinzeP.I.) with results presented byDelbo et al. (2003). For thethird piece, the most extensive NEO discovery statisare provided by the LINEAR survey(Stokes et al., 2002Stuart, 2001). The work ofStuart (2003)brings the fourthpiece, which appears here as a companion paper byStuartand Binzel (2004).

The observations presented here were obtained asof the Small Main-Belt Asteroid Spectroscopic Surv(SMASS) initiated at the Massachusetts Institute of Techogy in the early 1990’s. SMASS was undertaken to extour knowledge of the compositional properties of main-basteroids to smaller and smaller sizes by taking advanof state of the art charged–coupled device (CCD) detecThe first stage of this survey, which we herein refer toSMASSI, is reported byXu et al. (1995). The second stageSMASSII, is reported byBus and Binzel (2002a). Near-Earth objects, which we define as belonging to the AApollo, and Amor groups generally having perihelion dtances less than 1.3 AU from the Sun(Shoemaker et al1979), and Mars-crossing objects have been routinelyserved as targets of opportunity within the SMASS progsince its inception. As the SMASSII program achievedgoals (more than 1300 main-belt asteroids measured anextended system of asteroid taxonomy;Bus, 1999; Bus andBinzel, 2002a, 2002b), increasing focus has been directtoward the measurement of NEOs. This focus has been eabled by the increasing discovery rate and availabilityNEO targets for measurement(Stokes et al., 2002).

The fundamental scientific rationale for samplingphysical properties of near-Earth objects is reviewedBinzel et al. (2002). NEOs are “immigrants” to the inner Solar System with lifetimes of order 106–107 years(Morbidelliet al., 2002a). Because these lifetimes are extremely shcompared to the age of the Solar System, the current potion must be re-supplied. We seek to understand the oof the NEO population from both asteroidal and posscometary sources. What’s more, meteorites are (by dition of their orbital intersection) near-Earth objects prto their arrival—thus correlations between NEOs andteorites represent the most direct link for achieving anderstanding of asteroid–meteorite relationships. In addiNEOs are the smallest telescopically observable Solartem bodies. By using NEOs to leverage over the greapossible size range, we seek to investigate possible cositional trends as a function of size that may be indicativ

-

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.

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processes such as space weathering that modify the reflective properties of asteroid surfaces.

By virtue of their proximity, NEOs are among the moaccessible spacecraft destinations in our Solar System.ple return missions (e.g.,Sears et al., 2000) have the po-tential to address these science objectives (and more)the full capabilities of ground based analysis facilities. Rconnaissance of NEOs to select the most scientificallywarding candidates for sample return missions is a furmotivation for the SMASS NEO observations(Binzel et al.,2004b). Moving from the scientific to the pragmatic, asseing the NEO impact hazard(Morrison et al., 2002)requiresknowledge of the compositional distribution of the poplation. From the compositional distribution, albedo modcan be applied to convert the discovery H magnitude dibution into an actual size distribution(Delbo et al., 2003Stuart, 2003; Stuart and Binzel, 2004). Similarly, density as-sumptions can be applied to taxonomic categories to ymass estimates for impact energy distributions.Binzel et al.(2003)argue that at this stage of our knowledge, the msurement goals for scientifically understanding the NEOsthe same as the measurement goals for improving our asment of their impact hazard.

We organize this paper as follows. InSection 2we presena description and compilation of the SMASS NEO obsertions with the newly reported spectra being displayed inAppendix.Section 3presents a taxonomic analysis of tSMASS NEO data as well as other published NEO resThroughout this paper we utilize theBus (1999)taxonomy,with the addition of albedo data (where available) to disguish E-, M-, and P-types. These taxonomic results probasic input to the debiased population model derived incompanion paper(Stuart and Binzel, 2004). Section 4looksat the taxonomic and dynamic traceability of NEOs to thsource regions whileSection 5reveals new trends that mabe related to space weathering.Section 6presents a discussion of these trends and implications for the comecontribution to the NEO population. Concluding remarksgiven inSection 7.

2. SMASS observations

Table 1provides a summary compilation of SMASS rsults for 401 near-Earth and Mars-crossing objects. Throthe application of similar observing and reduction produres on a consistent set of telescopes, this table providelargest available uniform data set for this population. Witthe table, MC, ATE, APO, AMO denotes the orbital claas Mars-crossing, Aten, Apollo, or Amor. The taxonomtype, Slope, and principal component (PC) scores folthe system ofBus (1999), as discussed inSection 3. Theabsolute (H) magnitudes are from the Minor Planet Cedatabase. The final column ofTable 1includes referencefor the objects whose SMASS spectra have been previopublished. New SMASS spectra for 252 of these entries

Spectral properties of near-Earth objects 261

Table 1SMASS observational results for near-Earth and Mars-crossing objects

Number and name Provisional designation Orbit Type Slope PC2′ PC3′ H Mag T Reference

132 Aethra 1953 LF MC Xe 0.1605 0.2078 0.0063 9.4 3.18 2391 Ingeborg 1934 AJ MC S 0.5312 −0.1769 0.0150 10.1 3.41 2433 Eros 1898 DQ AMO S 0.6271 −0.2581 −0.0048 11.2 4.58 ∗512 Taurinensis 1903 LV MC S 0.4488 −0.2840 −0.0001 10.7 3.62 2699 Hela 1957 WX1 MC Sq 0.2152 −0.0519 0.0129 11.7 3.24 2719 Albert 1911 MT AMO S 0.4553 −0.1187 −0.0205 15.8 3.14 ∗985 Rosina 1922 MO MC S 0.5062 −0.2221 −0.0030 12.7 3.54 2

1011 Laodamia 1924 PK MC S 0.5650 −0.2118 −0.0421 12.7 3.44 ∗1036 Ganymed 1924 TD AMO S 0.4298 −0.2054 0.0474 9.5 3.03 ∗1065 Amundsenia 1955 SM1 MC S 0.5563 −0.2377 0.0291 13.2 3.48 21131 Porzia 1929 RO MC S 0.5324 −0.2669 −0.0055 13.0 3.59 21134 Kepler 1951 SA MC S 0.5075 −0.1024 0.0042 14.3 3.17 21139 Atami 1929 XE MC S 0.4885 −0.1872 −0.0021 12.5 3.82 21198 Atlantis 1931 RA MC L 0.7604 0.0984 0.0827 14.6 3.55 11204 Renzia 1931 TE MC S 0.4704 −0.1047 −0.0037 12.2 3.56 21293 Sonja 1933 SO MC Sq 0.2298 −0.2789 0.0069 12.0 3.59 21316 Kasan 1978 WK14 MC Sr 0.2799 −0.3933 −0.0305 13.3 3.34 21374 Isora 1935 UA MC Sq 0.2313 −0.1748 −0.0445 13.5 3.57 21565 Lemaitre 1948 WA MC Sq 0.1831 −0.2210 −0.0269 12.3 3.36 21593 Fagnes 1951 LB MC S 0.5191 −0.3458 −0.0317 13.2 3.57 21620 Geographos 1951 RA APO S 0.3892 −0.1177 0.0190 16.5 5.07 ∗1627 Ivar 1929 SH AMO S 0.5616 −0.3029 −0.0515 13.2 3.88 ∗1640 Nemo 1951 QA MC S 0.4981 −0.0457 0.0028 13.1 3.51 21660 Wood 1953 GA MC S 0.4426 −0.1498 0.0310 11.9 3.38 21685 Toro 1948 OA APO S 0.3031 −0.2011 −0.0146 14.0 4.71 61862 Apollo 1932 HA APO Q 0.0658 −0.3680 −0.0251 16.3 4.41 ∗1863 Antinous 1948 EA APO Sq 0.1776 −0.1793 0.0521 15.8 3.30 31864 Daedalus 1971 FA APO Sr 0.2560 −0.3072 −0.0780 15.0 4.33 ∗1865 Cerberus 1971 UA APO S 0.3085 −0.1508 −0.0047 17.0 5.59 ∗1866 Sisyphus 1972 XA APO S 0.6488 −0.1553 −0.0302 13.0 3.51 ∗1916 Boreas 1953 RA AMO S 0.5044 −0.1420 −0.0168 15.0 3.44 ∗1917 Cuyo 1968 AA AMO Sl 0.7233 −0.2161 0.0105 13.9 3.43 ∗1943 Anteros 1973 EC AMO L 0.6400 −0.0160 −0.0690 16.0 4.64 61951 Lick 1949 OA MC A 1.1774 −0.5920 −0.0676 14.7 4.54 21980 Tezcatlipoca 1950 LA AMO Sl 0.6492 −0.2832 −0.0045 14.0 3.99 ∗1981 Midas 1973 EA APO V −0.3096 −0.7383 0.0470 15.2 3.61 32035 Stearns 1973 SC MC Xe 0.4441 0.1200 0.0128 12.6 3.82 22062 Aten 1976 AA ATE Sr 0.3008 −0.3810 −0.0785 17.1 6.18 ∗2063 Bacchus 1977 HB APO Sq 0.2484 −0.2110 −0.0685 17.1 5.67 52064 Thomsen 1942 RQ MC S 0.4810 −0.1581 −0.0701 13.1 3.60 22074 Shoemaker 1974 UA MC Sa 0.8070 −0.1836 0.0649 14.0 3.90 12078 Nanking 1975 AD MC Sq 0.2431 −0.1220 0.0117 12.1 3.37 ∗2099 Opik 1977 VB MC Ch −0.0292 0.3217 −0.0673 15.2 3.36 22100 Ra-Shalom 1978 RA ATE C 0.0765 0.1810 0.0803 16.1 6.94 ∗2102 Tantalus 1975 YA APO Q −0.0186 −0.3202 −0.0192 16.2 4.45 ∗2201 Oljato 1947 XC APO Sq −0.0055 −0.1095 0.0441 16.9 3.30 ∗2204 Lyyli 1943 EQ MC X 0.3596 0.3154 0.0434 12.7 3.22 12253 Espinette 1932 PB MC Sl 0.7030 −0.0634 0.0603 12.9 3.55 12335 James 1974 UB MC Sa 0.6903 −0.3627 −0.0020 13.8 3.41 ∗2340 Hathor 1976 UA ATE Sq 0.0888 −0.0527 0.0287 19.2 6.88 ∗2423 Ibarruri 1972 NC MC A 0.8595 −0.3971 −0.0363 13.2 3.62 ∗2629 Rudra 1980 RB1 MC B −0.2512 0.3722 0.1447 14.5 4.02 22744 Birgitta 1975 RB MC S 0.3820 −0.2635 −0.0435 14.8 3.51 23040 Kozai 1979 BA MC S 0.5770 −0.3017 −0.0152 14.5 3.63 23102 Krok 1981 QA AMO S 0.3911 −0.3109 −0.0513 15.6 3.55 ∗3103 Eger 1982 BB APO Xe 0.5988 0.0946 −0.0371 15.4 4.61 ∗3122 Florence 1981 ET3 AMO S 0.3136 −0.2703 0.0243 14.2 3.92 ∗3198 Wallonia 1981 YH1 MC S 0.4586 −0.0100 0.0954 12.3 3.58 23199 Nefertiti 1982 RA AMO Sq 0.1479 −0.0205 −0.0353 15.1 4.19 ∗3200 Phaethon 1983 TB APO B −0.1941 0.3243 0.0460 14.3 4.51 ∗3216 Harrington 1980 RB MC S 0.4936 −0.1840 0.0571 14.0 3.46 23255 Tholen 1980 RA MC S 0.3313 −0.2011 0.0300 13.6 3.36 2

(continued on next page)


Table 1 (continued)


3287 Olmstead 1981 DK1 MC L 0.6200 0.0130 −0.0081 15.0 3.46 23288 Seleucus 1982 DV AMO K 0.4974 −0.0240 −0.0790 15.3 3.66 ∗3352 McAuliffe 1981 CW AMO A 0.8718 −0.3484 −0.0783 15.8 3.88 ∗3401 Vanphilos 1981 PA MC S 0.4994 −0.1016 0.0296 12.6 3.37 23416 Dorrit 1931 VP MC Sa 0.8073 −0.3391 −0.0620 13.7 3.81 23443 Leetsungdao 1979 SB1 MC T 0.5491 0.1168 −0.0111 13.3 3.43 23552 Don Quixote 1983 SA AMO D 1.1147 0.4581 0.0703 13.0 2.31 ∗3581 Alvarez 1985 HC MC B −0.3237 0.3113 0.0375 12.1 3.04 23635 1981 WO1 MC S 0.4567 −0.2039 0.0788 14.8 4.00 23671 Dionysus 1984 KD AMO Cb −0.0630 0.2376 0.0444 16.7 3.43 ∗3674 Erbisbuhl 1963 RH MC Sk 0.2658 −0.1008 0.0235 12.1 3.37 ∗3691 Bede 1982 FT AMO Xc 0.2058 0.0798 −0.0490 14.9 3.98 ∗3737 Beckman 1983 PA MC S 0.5326 −0.1470 0.0071 13.0 3.33 23753 Cruithne 1986 TO ATE Q 0.0130 −0.2254 0.0731 15.1 5.92 ∗3800 Karayusuf 1984 AB MC S 0.5372 −0.3416 −0.0440 15.4 4.36 23833 Calingasta 1971 SC MC Cb −0.0195 0.2452 0.0406 15.0 3.54 23858 Dorchester 1986 TG MC Sa 0.6807 −0.5277 −0.1248 13.7 3.62 23873 Roddy 1984 WB MC S 0.4203 −0.0952 0.0663 12.0 3.85 23908 Nyx 1980 PA AMO V −0.5268 −0.9300 0.0429 17.4 3.78 ∗3920 Aubignan 1948 WF MC Sa 0.6048 −0.3837 −0.0658 13.2 3.56 24034 1986 PA APO O −0.1440 −0.3152 −0.0801 18.1 5.70 ∗4055 Magellan 1985 DO2 AMO V −0.2210 −1.1119 0.0336 14.8 3.88 ∗4142 Dersu-Uzala 1981 KE MC A 1.0608 −0.2850 −0.0009 13.6 3.79 64179 Toutatis 1989 AC APO Sk 0.2733 −0.0239 0.0074 15.3 3.15 ∗4183 Cuno 1959 LM APO Sq 0.0019 −0.1273 −0.0165 14.4 3.57 ∗4197 1982 TA APO Sq 0.1727 −0.0500 −0.0245 14.9 3.09 ∗4205 David Hughes 1985 YP MC Xe 0.2392 0.1707 0.0168 14.7 4.10 24222 Nancita 1988 EK1 MC S 0.3205 −0.2770 −0.0435 12.4 3.48 24276 Clifford 1981 XA MC Cb 0.0977 0.3230 0.1000 14.3 3.72 24341 Poseidon 1987 KF APO O −0.1709 −0.1810 −0.0576 15.5 3.69 ∗4435 Holt 1983 AG2 MC S 0.4085 −0.2522 0.0096 13.2 3.41 24451 Grieve 1988 JJ MC S 0.5270 −0.1972 0.0169 12.2 3.15 ∗4503 Cleobulus 1989 WM AMO Sq 0.1184 −0.1224 0.0049 15.6 3.15 ∗4558 Janesick 1988 NF MC S 0.3118 −0.2297 0.0299 12.2 3.49 24660 Nereus 1982 DB APO Xe 0.2958 0.2634 0.0409 18.3 4.49 64688 1980 WF AMO V 0.0852 −0.4530 −0.0582 19.0 3.44 ∗4910 Kawasato 1953 PR MC S 0.5497 −0.2466 −0.0053 14.2 3.42 24947 Ninkasi 1988 TJ1 AMO Sq 0.1981 −0.2952 −0.0278 18.7 4.77 ∗4954 Eric 1990 SQ AMO S 0.5589 −0.2375 0.0544 12.6 3.66 ∗4957 Brucemurray 1990 XJ AMO S 0.6394 −0.0886 −0.0466 15.1 4.20 ∗4995 1984 QR MC S 0.3525 −0.1473 0.0105 13.0 3.41 25038 Overbeek 1948 KF MC S 0.4931 −0.1058 −0.0213 14.1 3.51 25131 1990 BG APO S 0.5089 −0.1387 0.0031 14.1 4.21 ∗5143 Heracles 1991 VL APO O −0.1949 −0.1147 −0.0371 14.0 3.58 ∗5230 Asahina 1988 EF MC S 0.4305 −0.2170 −0.0127 13.4 3.35 25253 1985 XB MC S 0.6193 −0.2635 −0.0329 13.7 3.69 25275 Zdislava 1986 UU MC Sa 0.6226 −0.4940 −0.0639 13.6 3.62 ∗5349 Paulharris 1988 RA MC C −0.0416 0.2833 0.0030 12.7 3.00 25392 Parker 1986 AK MC Sl 0.6551 −0.2174 −0.0876 12.7 3.38 25407 1992 AX MC Sk 0.2271 −0.0085 0.0099 13.7 3.95 25510 1988 RF7 MC S 0.6165 −0.2987 −0.0099 13.9 3.63 25585 Parks 1990 MJ MC Ch −0.0644 0.3847 0.0111 13.7 3.08 25587 1990 SB AMO Sq 0.2478 −0.0442 0.3496 13.6 3.25 ∗5604 1992 FE ATE V −0.1308 −1.2293 0.0065 16.4 6.38 ∗5626 1991 FE AMO S 0.5842 −0.2294 0.0263 14.7 3.52 ∗5641 McCleese 1990 DJ MC A 0.8533 −0.3894 −0.0773 12.7 3.94 ∗5646 1990 TR AMO U −0.1015 −0.6302 −0.0859 16.1 3.57 ∗5649 Donnashirley 1990 WZ2 MC S 0.3343 −0.2600 0.0041 13.2 3.44 25660 1974 MA APO Q 0.1450 −0.3748 −0.1204 15.7 3.51 ∗5732 1988 WC MC S 0.4511 −0.1602 0.0177 14.1 3.45 25751 Zao 1992 AC AMO X 0.2418 0.1032 0.0272 14.9 3.58 ∗5817 Robertfrazer 1989 RZ MC S 0.5106 −0.3254 0.0172 12.6 3.35 25828 1991 AM APO Q −0.0250 −0.1546 −0.0113 16.3 3.77 ∗



Table 1 (continued)


5836 1993 MF AMO S 0.4099 −0.1464 0.0222 15.0 3.28 ∗6047 1991 TB1 APO S 0.5319 −0.1278 −0.0136 17.0 4.48 ∗6053 1993 BW3 AMO Sq 0.2602 −0.2218 0.0279 15.2 3.44 ∗6249 Jennifer 1991 JF1 MC Xe 0.4044 0.1247 0.0499 12.4 3.78 26386 1989 NK1 MC S 0.5220 −0.1874 0.0601 12.7 3.54 26455 1992 HE APO S 0.4823 −0.3305 0.1027 13.8 3.18 ∗6489 Golevka 1991 JX APO Q 0.2310 −0.3507 −0.0482 19.1 3.18 ∗6500 Kodaira 1993 ET MC B −0.2539 0.3066 0.0620 12.5 3.04 26569 1993 MO AMO Sr 0.4724 −0.4910 −0.1319 16.2 4.21 ∗6585 O’Keefe 1984 SR MC Sk 0.2748 −0.1441 −0.0192 14.3 3.36 26611 1993 VW APO V −0.1473 −0.7528 0.1062 16.5 4.05 ∗6847 Kunz-Hallstei 1977 RL MC Sk 0.2590 −0.0811 −0.0009 13.7 3.40 27304 Namiki 1994 AE2 MC Ld 0.8737 0.0157 0.1074 13.3 3.25 27336 Saunders 1989 RS1 AMO Sq −0.0671 −0.1819 0.0018 18.7 3.41 ∗7341 1991 VK APO Sq 0.0632 −0.2138 −0.0211 16.7 3.84 ∗7358 Oze 1995 YA3 AMO Sq 0.1666 −0.1383 0.0000 14.4 3.49 ∗7474 1992 TC AMO X 0.3486 0.3191 0.0702 18.0 4.36 17480 Norwan 1994 PC AMO S 0.3011 −0.2670 0.0060 17.5 4.34 ∗7482 1994 PC1 APO S 0.4594 −0.2495 −0.0101 16.8 4.66 ∗7604 1995 QY2 MC C −0.0702 0.1381 0.0179 13.7 2.86 27753 1988 XB APO B −0.2817 0.3549 0.1009 18.6 4.47 57822 1991 CS APO S 0.4548 −0.1713 −0.0176 17.4 5.36 ∗7888 1993 UC APO U 1.0453 −0.7483 −0.2914 15.3 3.05 ∗7889 1994 LX APO V −0.4210 −1.2968 0.1178 15.3 4.86 ∗7977 1977 QQ5 AMO S 0.3087 −0.1537 0.0584 15.4 3.38 ∗8176 1991 WA APO Q 0.0078 −0.2064 −0.0218 17.1 3.95 ∗8201 1994 AH2 APO O −0.1989 −0.2479 0.0195 16.3 3.02 38566 1996 EN APO U −0.2108 −1.6055 0.1226 16.5 4.22 ∗9400 1994 TW1 AMO Sr 0.4645 −0.4771 −0.0901 14.8 2.94 ∗9969 Braille 1992 KD MC Q 0.0863 −0.2935 −0.0896 15.8 3.28 3

10115 1992 SK APO S: 17.0 5.06 ∗10165 1995 BL2 APO L 0.6313 −0.0295 −0.0181 17.1 4.98 ∗10302 1989 ML AMO X 0.1365 0.2478 0.0443 19.5 5.06 310563 Izhdubar 1993 WD APO Q 0.0549 −0.1667 −0.0073 17.3 5.54 ∗11066 Sigurd 1992 CC1 APO K 0.4980 0.0917 0.0992 15.0 4.50 111311 Peleus 1993 XN2 APO Sq 0.2460 −0.1192 0.0050 16.5 3.43 ∗11398 1998 YP11 AMO Sr 0.2792 −0.3033 −0.0382 16.3 4.05 ∗11405 1999 CV3 APO Sq 0.2284 −0.2422 −0.0182 15.0 4.46 ∗11500 1989 UR APO S 0.3906 −0.1247 −0.0103 18.4 5.65 ∗12538 1998 OH APO S: 16.1 4.28 ∗12711 1991 BB APO Sr 0.3301 −0.3537 −0.0808 16.0 5.10 ∗12923 1999 GK4 APO S: 16.1 3.72 ∗13651 1997 BR APO S 0.3346 −0.2565 −0.1260 17.6 4.81 ∗14402 1991 DB AMO C 0.0077 0.2024 0.0214 18.4 4.06 ∗15745 1991 PM5 AMO S 0.5345 −0.1150 0.0078 17.8 4.10 ∗15817 Lucianotesi 1994 QC AMO Xc 0.1489 0.0752 −0.1092 18.6 4.90 ∗16064 1999 RH27 AMO C 0.1197 0.2540 0.0044 16.9 3.02 316657 1993 UB AMO Sr 0.2720 −0.2628 0.0622 16.9 3.35 ∗16834 1997 WU22 APO S 0.5482 −0.0848 −0.0189 15.7 4.46 316960 1998 QS52 APO Sq −0.0205 −0.0503 −0.0139 14.3 3.00 ∗17274 2000 LC16 AMO Xk 0.3726 0.2333 0.0021 16.7 3.10 ∗17511 1992 QN APO X 0.1482 0.2297 0.0308 17.1 5.25 ∗18514 1996 TE11 MC Xc 0.1238 0.1468 −0.0007 15.8 3.15 218736 1998 NU AMO Sk 0.2540 −0.0260 0.0737 16.1 3.38 ∗18882 1999 YN4 AMO S 0.3034 −0.1712 0.0164 16.3 3.97 319356 1997 GH3 AMO S 0.3065 −0.2516 −0.0467 17.1 3.22 ∗20043 1993 EM MC U −0.1013 0.0583 0.0576 15.4 3.81 320255 1998 FX2 AMO Sq 0.2655 −0.1376 0.0101 18.2 3.52 ∗20425 1998 VD35 APO Sq 0.1093 −0.3739 −0.0595 20.4 4.28 ∗20790 2000 SE45 AMO S 0.3174 −0.0482 −0.0092 16.6 3.09 ∗20826 2000 UV13 APO Sq 0.2232 −0.1018 0.0049 13.5 3.04 ∗22099 2000 EX106 APO S: 18.0 5.58 ∗22449 1996 VC MC S 0.3817 −0.1729 −0.0513 13.8 3.39 2



Table 1 (continued)


22771 1999 CU3 APO Sl 0.6226 −0.0570 −0.0364 17.0 4.22 ∗23548 1994 EF2 AMO Q 0.1469 −0.3157 0.0033 17.6 3.31 ∗24475 2000 VN2 AMO Sa 0.7374 −0.2446 −0.0034 16.5 3.70 ∗25143 Itokawa 1998 SF36 APO S(IV) 0.4644 −0.3073 −0.2203 19.2 4.90 425330 1999 KV4 APO B −0.0545 0.2652 0.0129 16.8 4.35 ∗26209 1997 RD1 MC Sq 0.2589 −0.1600 −0.0158 16.0 3.23 231345 1998 PG AMO Sq 0.1485 −0.2232 −0.0576 17.6 3.72 ∗31346 1998 PB1 AMO Sq: 17.1 3.68 ∗32906 1994 RH AMO S 0.2340 −0.0092 0.0012 16.0 3.43 ∗35107 1991 VH APO Sk 0.3109 −0.1960 −0.1196 16.9 5.47 ∗35396 1997 XF11 APO Xk 0.2616 0.0356 −0.0273 16.9 4.53 535432 1998 BG9 AMO S 0.3359 −0.1532 0.0295 19.5 3.21 335670 1998 SU27 APO Sq 0.0195 −0.0845 −0.0129 19.6 3.47 ∗36017 1999 ND43 AMO Sl 0.6535 −0.1569 0.0989 19.2 4.44 336183 1999 TX16 AMO Ld 0.9304 −0.0633 0.0375 16.2 4.16 336284 2000 DM8 APO Sq 0.2411 −0.1363 −0.0095 14.9 4.11 ∗37336 2001 RM AMO S 0.2976 −0.0851 −0.0117 15.2 3.23 ∗40267 1999 GJ4 APO Sq 0.0082 −0.2570 0.0999 15.0 4.38 340310 1999 KU4 MC S: 16.8 3.55 ∗47035 1998 WS MC Sr 0.2595 −0.2472 0.0605 12.5 3.12 248603 1995 BC2 AMO X 0.1629 0.1255 0.0213 17.3 3.80 ∗53319 1999 JM8 APO X: 15.2 2.99 ∗

1989 UQ ATE B −0.0991 0.2898 0.0625 19.0 6.49 51989 VA ATE Sq −0.0155 −0.2709 −0.0440 17.9 7.66 ∗1991 BN APO Q −0.0125 −0.1486 −0.0345 19.3 4.57 ∗1991 XB AMO K 0.3583 0.1423 0.0745 18.1 2.93 11992 BF ATE Xc 0.1581 0.0422 −0.0434 19.5 6.52 ∗1992 NA AMO C −0.0687 0.3189 0.0517 16.5 3.28 11992 UB AMO X 0.3485 0.2866 0.0941 16.0 2.90 11993 TQ2 AMO Sa 0.6672 −0.2869 0.0048 20.0 3.73 ∗1994 AB1 AMO Sq 0.2612 −0.1357 0.0421 16.3 3.02 ∗1994 AW1 AMO L 0.6599 0.0345 −0.0630 17.7 5.55 ∗1994 TF2 ATE Sr 0.3730 −0.3147 −0.1242 19.3 6.00 ∗1995 BM2 MC Sq 0.2142 −0.2681 0.0153 15.2 3.43 21995 WL8 AMO Sq 0.0950 −0.0616 0.0655 18.1 3.32 ∗1995 WQ5 MC Ch −0.1551 0.3241 0.0139 16.8 3.36 21996 BZ3 AMO X 0.1665 0.1918 0.0363 18.2 3.18 ∗1996 FG3 APO C −0.0786 0.2497 −0.0122 17.8 5.78 31996 FQ3 AMO Sq 0.2013 −0.1803 0.0854 21.0 3.66 ∗1996 GT AMO Xk 0.4066 0.1959 −0.0164 18.5 4.20 51996 UK MC Sq 0.1050 −0.1322 0.0333 16.4 3.19 21997 AC11 ATE Xc 0.0658 0.1867 0.0106 21.0 6.36 ∗1997 AQ18 APO C −0.0333 0.2344 0.0185 18.2 5.33 ∗1997 BQ APO S 0.6060 −0.2500 −0.0587 18.0 3.98 ∗1997 CZ5 MC S 0.3600 −0.2067 0.0000 13.6 3.37 ∗1997 GL3 APO V −0.3161 −0.4887 0.1259 20.0 3.10 ∗1997 RT AMO O −0.1940 −0.2061 0.0186 20.0 3.43 ∗1997 SE5 AMO T 0.6256 0.2200 0.0866 14.8 2.66 ∗1997 TT25 AMO Sq 0.2395 −0.1453 0.0063 19.3 3.60 ∗1997 UH9 ATE Sq 0.1526 −0.1328 −0.0146 18.8 6.90 ∗1997 US9 APO Q −0.0392 −0.1906 −0.0305 17.3 5.75 ∗1998 BB10 APO Sq 0.1320 −0.1429 0.0808 20.4 4.97 31998 BM10 MC Sq 0.2596 −0.1225 0.0496 16.5 3.31 31998 BT13 APO Sq 0.1269 −0.1703 0.0592 26.5 3.19 31998 FM5 AMO S 0.4468 −0.2091 0.0017 16.0 3.37 ∗1998 HT31 APO C: 20.8 3.05 ∗1998 KU2 AMO Cb −0.0044 0.2611 0.0877 16.6 3.40 ∗1998 MQ AMO S 0.3283 −0.3191 0.0127 16.6 3.89 ∗1998 MW5 APO Sq 0.0694 −0.3656 −0.0544 19.2 5.90 31998 QR15 AMO Sq 0.0048 −0.1452 −0.0547 18.0 3.07 ∗1998 SG2 AMO Sq −0.0695 −0.1749 −0.0535 19.7 3.48 ∗1998 ST49 APO Q −0.0904 −0.2829 0.1064 17.7 3.23 31998 UT18 APO C −0.0964 0.1721 −0.0078 20.4 3.01 3



Table 1 (continued)


1998 VO33 APO V −0.2417 −0.3718 0.0687 17.0 4.72 ∗1998 VR ATE Sk 0.2964 −0.1185 −0.0574 18.5 6.66 ∗1998 WM APO Sq 0.1664 −0.2466 −0.0306 16.8 5.10 ∗1998 WP5 AMO Sl 0.7720 −0.2027 0.0391 18.4 4.74 31998 WZ6 APO V −0.1423 −0.4932 0.1462 17.3 4.46 ∗1998 XB ATE S: 15.5 6.49 ∗1998 XM4 APO S 0.5354 −0.1596 −0.1655 15.4 3.61 ∗1999 AQ10 ATE S: 20.3 6.37 ∗1999 CF9 APO Q −0.1338 −0.1323 0.0809 17.8 3.86 31999 CV8 APO V −0.0074 −0.5133 −0.0277 19.6 4.91 ∗1999 CW8 APO B −0.1459 0.2284 0.0296 18.5 3.20 ∗1999 DB2 AMO Sq 0.2182 −0.2762 −0.0697 19.1 2.90 ∗1999 DJ4 APO Sq 0.1124 −0.2085 −0.0165 18.5 3.84 31999 DY2 AMO Sr 0.2861 −0.3287 −0.0335 21.9 3.65 ∗1999 EE5 AMO S 0.3304 −0.2154 −0.0201 18.4 4.05 ∗1999 FA APO S 0.5267 −0.1734 −0.0597 20.5 5.70 51999 FB APO Q −0.0370 −0.1015 0.0264 18.1 5.14 ∗1999 FK21 ATE S 0.3880 −0.1632 0.0481 18.9 7.56 ∗1999 FN19 AMO Sq 0.2164 −0.1613 −0.0064 22.5 4.19 51999 HF1 ATE X: 14.5 6.98 ∗1999 JD6 ATE K 0.5027 0.0223 −0.1584 17.2 6.50 31999 JE1 APO Sq −0.0829 −0.2265 0.0664 19.5 4.59 31999 JO8 AMO S 0.3277 −0.2690 −0.0018 17.0 3.02 31999 JU3 APO Cg 0.0442 0.1003 −0.1112 19.6 5.31 31999 JV3 APO S 0.5743 −0.2689 −0.0505 19.0 4.51 31999 JV6 APO Xk 0.2799 0.0553 −0.1780 19.9 6.00 31999 KW4 ATE S: 16.6 8.50 ∗1999 NC43 APO Q 0.0140 −0.3559 0.0138 16.0 3.90 ∗1999 RB32 AMO V −0.0573 −0.5767 −0.0557 19.8 3.26 31999 SE10 AMO X 0.5135 0.2743 −0.0380 20.0 2.84 31999 SK10 APO Sq 0.1469 −0.1628 −0.0459 19.3 3.99 31999 VM40 AMO S 0.5702 −0.1543 −0.0136 14.6 3.37 31999 VN6 AMO C 0.0574 0.2430 0.0267 19.5 4.01 31999 VQ5 MC Q −0.1000 −0.2316 −0.0248 19.4 3.13 31999 WK13 AMO S 0.4297 −0.1468 0.0080 17.2 3.74 31999 XO35 AMO Sq 0.1758 −0.2042 −0.0196 16.8 3.13 31999 YB AMO Sq 0.2730 −0.2093 −0.1006 18.5 4.93 31999 YD AMO Sk 0.2899 −0.1230 −0.1395 21.1 3.22 31999 YF3 AMO Sq 0.2502 −0.0966 −0.0468 18.5 4.44 31999 YG3 APO S 0.5264 −0.1113 0.0322 19.1 4.82 31999 YK5 ATE X 0.1050 0.2749 0.0732 16.8 6.91 32000 AC6 ATE Q −0.0917 −0.1754 −0.0144 21.0 6.87 32000 AE205 AMO S 0.5336 −0.1871 −0.0714 22.9 5.41 32000 AH205 APO Sk 0.3281 −0.2121 −0.1680 22.4 5.39 32000 AX93 AMO Sq 0.1494 −0.1547 0.0103 17.7 3.40 32000 BG19 AMO X 0.3114 0.1296 0.0425 17.9 3.11 ∗2000 BJ19 APO Q −0.0330 −0.1034 −0.0099 16.2 4.58 ∗2000 BM19 ATE O −0.0748 −0.2909 0.0817 18.2 7.72 52000 CE59 APO L 0.6373 0.0301 −0.0086 20.4 5.47 ∗2000 CK33 ATE Xk 0.3880 0.1004 −0.1426 18.2 6.12 ∗2000 CN33 AMO X 0.3049 0.2639 0.0131 19.2 3.09 ∗2000 CO101 APO Xk 0.2553 0.2072 −0.0039 19.3 5.71 ∗2000 DO1 APO V −0.1613 −0.5764 0.1665 20.4 4.41 ∗2000 DO8 APO S: 24.8 3.19 ∗2000 EA107 ATE Q 0.0135 −0.1665 0.0143 16.2 6.25 ∗2000 EZ148 APO S: 15.5 3.10 ∗2000 GD2 ATE Sq 0.1182 −0.2621 −0.1422 19.2 7.43 ∗2000 GJ147 APO S: 19.5 5.31 ∗2000 GK137 APO Sq 0.1612 −0.4228 0.0349 17.4 3.66 ∗2000 GO82 APO S: 16.8 3.09 ∗2000 GR146 APO S: 16.3 4.40 ∗2000 GU127 APO S: 18.7 3.33 ∗2000 GV127 AMO S: 19.2 2.94 ∗



Table 1 (continued)


2000 JG5 APO S: 18.3 4.40 ∗2000 JQ66 AMO R 0.1709 −0.5579 −0.0166 18.1 3.56 ∗2000 KL33 AMO S 0.4014 −0.2203 −0.0499 19.7 3.60 ∗2000 MU1 APO S 0.3325 −0.1983 −0.0468 19.9 4.71 ∗2000 NM APO Sr 0.2650 −0.2776 −0.0259 15.6 2.93 ∗2000 OJ8 AMO Sr 0.2645 −0.2812 −0.0603 16.8 3.31 ∗2000 PG3 APO D 0.8912 0.3498 0.0339 16.2 2.55 ∗2000 RW37 APO C 0.1129 0.1902 0.0312 20.2 5.09 ∗2000 SY162 AMO Sq: 0.1228 −0.0981 −0.1457 19.3 3.44 ∗2000 WC67 AMO X 0.3689 0.4058 −0.0624 19.1 3.09 ∗2000 WF6 AMO Sq 0.1118 −0.1327 −0.0243 18.7 3.04 ∗2000 WJ10 AMO Xk 0.3045 0.1022 −0.0856 20.6 3.59 ∗2000 WJ63 AMO Sq 0.0740 −0.1820 −0.0577 20.9 3.02 ∗2000 WK10 APO X 0.3869 0.4195 0.0015 18.5 4.25 ∗2000 WL10 APO Xc 0.2051 0.2908 −0.0089 18.0 2.72 ∗2000 WL63 AMO S 0.3429 0.0344 −0.0231 20.4 4.58 ∗2000 WM63 APO S 0.5154 −0.0725 −0.0570 20.2 5.87 ∗2000 WO107 ATE X 0.2668 0.3169 −0.0421 19.4 6.23 ∗2000 XL44 AMO S 0.3619 −0.0498 0.0778 18.0 3.50 ∗2000 YA APO Sk 0.2565 −0.0491 −0.0387 23.6 3.21 ∗2000 YF29 APO S 0.5054 −0.0409 0.0067 20.2 4.47 52000 YH66 APO Xk 0.3397 0.2441 −0.1048 17.5 5.04 ∗2000 YO29 APO C 0.0354 0.2222 −0.0539 18.0 3.36 ∗2001 AE2 AMO T 0.5668 0.2436 −0.1098 19.2 4.87 52001 CC21 APO L 0.4807 0.0529 −0.0668 18.4 5.91 52001 DU8 AMO S 0.5482 −0.0524 −0.0183 16.4 3.85 ∗2001 EB AMO Sl 0.5774 −0.0567 0.0056 17.3 4.07 ∗2001 EC APO Sq 0.1184 −0.1831 0.0232 18.6 2.91 ∗2001 FY AMO S 0.4578 −0.1350 −0.0223 18.9 3.89 ∗2001 HA8 AMO C: −0.1543 0.4436 0.0146 16.9 3.31 ∗2001 HK31 AMO X 0.3964 0.2757 −0.0072 21.0 3.25 ∗2001 HW15 AMO X 0.1685 0.2479 0.0028 20.2 4.40 ∗2001 JM1 AMO S 0.4320 −0.1749 −0.0220 19.0 4.52 ∗2001 JV1 APO Sq 0.1545 −0.1214 0.0050 21.3 4.08 ∗2001 MF1 AMO Sk 0.2617 −0.1227 0.0252 16.8 3.03 ∗2001 OE84 AMO S 0.3952 −0.1748 −0.0257 17.8 3.43 ∗2001 PD1 AMO K: 0.4288 0.0441 0.0519 18.2 3.49 ∗2001 QA143 AMO Sk 0.2827 −0.0403 −0.0250 19.6 3.44 ∗2001 QQ142 APO Sq 0.1837 0.0108 0.0629 18.5 4.64 ∗2001 SG10 APO X 0.2415 0.3184 0.0318 20.3 4.54 52001 SG286 APO D 1.2455 0.3639 −0.0237 21.1 4.77 52001 SJ262 AMO C: −0.0494 0.2849 −0.0085 19.6 2.98 ∗2001 SK162 AMO T 0.5451 0.2713 0.0104 17.9 3.77 52001 TC45 APO Sq 0.0674 0.0230 0.1085 19.1 3.31 ∗2001 TX16 MC X 0.3294 0.2983 0.0137 14.1 2.77 ∗2001 TY44 AMO X 0.3215 0.4383 −0.0058 20.3 3.35 ∗2001 UA5 APO Sq 0.2092 −0.0881 0.0642 17.4 3.94 ∗2001 UC5 AMO X 0.2903 0.4534 0.0275 21.3 2.87 ∗2001 UU92 AMO T 0.5228 0.3318 −0.1421 19.8 2.80 ∗2001 UY4 APO X 0.1920 0.2353 0.0021 18.4 4.23 ∗2001 VG5 APO Sq 0.2176 −0.1558 −0.0288 16.7 3.28 ∗2001 VS78 AMO S 0.3394 −0.2435 −0.0211 15.5 3.94 ∗2001 WA25 APO S 0.2973 −0.2064 −0.0618 18.7 3.90 ∗2001 WG2 APO Sk 0.2598 −0.0519 −0.0570 16.3 3.56 ∗2001 WH2 AMO X 0.2472 0.3780 −0.0078 20.0 3.65 ∗2001 WL15 AMO Sk 0.2623 −0.0668 −0.0806 18.3 3.56 ∗2001 XN254 AMO S 0.4937 −0.1993 0.0590 17.5 3.35 ∗2001 XR1 APO Sq 0.1650 −0.1272 0.0019 17.4 4.95 ∗2001 XS1 AMO Cb 0.1222 0.3537 0.0172 18.8 3.14 ∗2001 XS30 APO Xc 0.0369 0.1116 −0.0127 17.5 4.93 ∗2001 XU30 APO Sq 0.1909 −0.0616 −0.0501 19.9 3.33 ∗2001 XY10 ATE Sk 0.2047 0.0281 −0.0441 20.4 6.66 ∗2001 YE1 APO T 0.6117 0.2376 −0.0799 20.6 3.82 ∗



Table 1 (continued)


2001 YK4 APO X: 0.2210 0.2629 0.0195 18.5 2.83 ∗2002 AA APO Sq 0.2389 −0.1375 0.0363 19.5 5.42 ∗2002 AD9 APO L 0.6439 0.0097 −0.0828 16.5 3.58 ∗2002 AH29 AMO K 0.3581 0.0702 −0.0019 21.7 3.29 ∗2002 AK14 APO V: −0.0590 −0.4693 0.0365 21.7 5.97 ∗2002 AL14 ATE Ld 0.9826 0.1559 0.0825 17.8 6.26 ∗2002 AL31 APO X 0.2859 0.2359 −0.0202 24.4 5.34 52002 AQ2 AMO S 0.3978 −0.0367 −0.1564 18.6 3.08 ∗2002 AT4 AMO D 0.8578 0.2713 0.0104 20.9 3.95 52002 AU5 APO X 0.3663 0.2713 0.0104 17.7 3.56 ∗2002 AV APO K 0.2958 0.0557 0.0109 20.7 3.13 ∗2002 BA1 AMO S 0.4377 −0.2128 0.0042 21.7 3.73 ∗2002 BK25 APO Sk 0.1673 0.0427 −0.0306 18.1 3.13 ∗2002 BM26 APO X 0.4393 0.2702 0.0408 20.1 3.86 ∗2002 BP26 AMO X 0.2082 0.2083 0.0498 19.3 3.97 ∗2002 CS11 AMO X: 0.2506 0.2713 0.0104 21.6 3.70 ∗2002 CT46 AMO Sr 0.3123 −0.3365 0.0107 20.9 3.30 ∗2002 DH2 APO Ch 0.0228 0.3159 0.0128 20.3 3.52 ∗2002 DO3 APO X: 0.2489 0.5103 0.0566 22.0 3.83 ∗2002 DQ3 AMO Sq 0.0135 0.0272 −0.0132 23.8 4.74 52002 DU3 APO Sq 0.2518 −0.2181 0.0157 20.8 5.44 52002 DY3 AMO Xk 0.2668 0.1899 0.0928 18.6 4.43 ∗2002 EA APO L 0.4936 0.0218 −0.0345 22.4 4.69 ∗2002 EC AMO X: 0.2545 0.3439 −0.2688 23.3 3.51 ∗

∗—This work; 1—Xu et al. (1995); 2—Bus and Binzel (2002a); 3—Binzel et al. (2001a); 4—Binzel et al. (2001b); 5—Binzel et al. (2004b); 6—Binzel et al.
(2004a).
thewly

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presented here, where these objects are denoted by “*” inReference column. Observing circumstances for all nereported measurements are presented inAppendix Aand thenewly reported spectra appear inAppendix Band are avail-able in digital format athttp://smass.mit.edu/.

For SMASS observations prior to 1997, visible wavlength measurements were made almost exclusively withMark III long-slit spectrograph coupled to the MichigaDartmouth-MIT Observatory (MDM) 2.4-m telescope lcated on the southwest ridge of Kitt Peak, Arizona. Oobserving procedures carried out there (fully describeXu et al., 1995; Bus, 1999; Bus and Binzel, 2002a) werefollowed closely at all sites. Particular details in comminclude using a long slit oriented north–south to simultaously image the object spectrum and the backgroundwith a slit width several times wider than the seeing diskthe best possible photometric precision. Comparable speresolution, typicallyλ/�λ ∼ 100 was obtained with all telescope/instrument/slit/detectorconfigurations. Spectral expo-sures, typically 900 s or shorter (to minimize the accumution of cosmic-ray hits on the detector) were nearly alwmade while the object was within one hour of the merian to minimize any effects of atmospheric dispersion.all observations we utilized Hyades 64 and 16 Cyg B asprimary reference stars for the solar analog spectrum.further sky coverage, we also utilized solar analog referestars selected fromLandolt (1973)that were verified to bewithin 1% of our primary reference stars, as detailed inBusand Binzel (2002a).

In 1998 a new collaboration began providing to tSMASS program visible wavelength measurements m

l

using the Double Spectrograph on the Palomar Obsetory 5-m Hale telescope (co-author AWH, Palomar PrinciInvestigator). Our use of this telescope and instrument cbination is fully described inBinzel et al. (2001a). SMASSNEO visible wavelength measurements were also beguing the RCSP spectrograph on the Kitt Peak Nationalservatory 4-m Mayall telescope (Binzel, P.I.) in 1999,detailed inBinzel et al. (2001b). While our CCD spectra obtained at MDM, Palomar, and Kitt Peak generally rangonly out to 0.92-µm, a program to extend measurementsto 1.6-µm was begun at the NASA Infrared Telescopecility (IRTF) in 1997. This so-called “SMASSIR” surveutilized a low-resolution “asteroid grism” system designby one of us (RPB) and described in Fig. 2 ofBinzel etal. (2001a). The performance of the asteroid grism andmethods utilized for the acquisition, reduction, and calibration of these data are detailed inBurbine (2000) and Burbinand Binzel (2002).

In this paper, we also present results from a one nvisible spectroscopy run on the Keck II 10-m. This Ketelescope night was allocated to obtain spectral measments of 4660 Nereus in support of the NASA collabotion with the MUSES-C mission. (Nereus was the initiaplanned MUSES-C target.) We utilized LRIS (low resotion imaging spectrometer) to obtain spectra over the walength range of 0.5- to 1.0-µm with a 300 line/mm grating.The basic observing, reduction, and solar calibration pcedures described above were also applied to thesemeasurements. While Nereus was not available duringentire night, four other NEOs were measured as tarof opportunity. Finally, we also incorporate results fro



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the newly operational Magellan I 6.5-m telescope atCampanas, Chile. There we utilized the Boller and Chivspectrograph, also using a 300 line/mm grating to cover thewavelength range of 0.5- to 0.9-µm. Similarly, our standobserving, reduction, and solar calibration procedures wfollowed.

3. Taxonomic analysis

3.1. Taxonomic classification

Our first analysis step is to place the newly observedjects into the taxonomic system ofBus (1999). The Bustaxonomy is based on a uniform data set of 1341 mbelt asteroid spectra from the SMASSII survey measuwith the 2.4-m telescope and Mark III spectrograph atMDM Observatory(Bus, 1999; Bus and Binzel, 2002a). Asour NEO survey utilized a wide variety of telescopes aspectrograph combinations, we chose to re-observe a ssubset of main-belt objects for direct comparison andrelation with the SMASSII survey results. To quantify ocomparison, we calculate spectral slope values and prpal component scores using theBus (1999)slope definitionand eigenvectors. Sufficient data were obtained to compathe principal component scores for 19 SMASSII main-basteroids re-observed at Palomar and 24 re-observed aPeak. As the Kitt Peak data generally do not extend be0.50-µm, extrapolation was necessary to the 0.44-µm lolimit of the eigenvectors. This extrapolation was madelowing the curvature (or linearity) occurring in the 0.5-0.6-µm spectral range. Both Palomar and Kitt Peak sets sa slight but mutually consistent offset from the SMAS

ll

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slope and PCA scores. This is a purely empirical offset,haps due to slightly different spectrophotometric responof the various systems. As shown inFig. 1, this offset is toosmall to affect any taxonomic classification except inimmediate vicinity of the class boundaries where a natambiguity always exists(Bus, 1999). However because wseek to perform statistical trend analysis, we must accfor this offset in order to achieve a data set that is as intenally consistent as possible. Letting Slope, PC2′, PC3′ be thevalues on the system ofBus (1999)and letting Slope0, PC20,PC30 be the initially calculated values for the new Palomand Kitt Peak data reported here, the transformations toSMASSII system(Bus, 1999)are given by:

Slope= 0.0402+ 0.6942× Slope0,

PC2′ = −0.0610+ PC20,

PC3′ = −0.0240+ 0.5070× PC30.

All resulting values for Slope, PC2′, and PC3′ are presentein Table 1. We note that no transformations were applto the 106 NEOs observed during the SMASSII surveying the MDM 2.4-m telescope as these are formally onSMASSII system. (These SMASSII measurements aretinguished inAppendix Aas having been made using MD2.4-m telescope.) Similarly, 10 NEOs observed in SMAS(Xu et al., 1995, also using the MDM 2.4-m telescophave their principal componentscores recalculated with thSMASSII eigenvectors but with no transformation appliThe resulting distribution of Slope and Principal Componescores comparing the near-Earth and main-belt populaare plotted inFig. 1.

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Fig. 1. Principal component space within theBus (1999)taxonomy system designed for visible wavelengthCCD spectra. The 1341 SMASS main-belt asterodefining this system are depicted by their classification letters while 400 SMASS near-Earth and Mars-crossing objects are depicted in red. NEOs shgreaterdispersion in spectral properties, and most notably, span the once empty gap between the S- and Q-complexes(Binzel et al., 1996, 2001a). The green linedenotes the defined boundary between the C- and X-complexes, where objects close to this line have a natural ambiguity in their taxonomic assignmenD-type objects in the upper right are candidates for extinct comets. The horizontal line at lowerright represents the magnitude of the average transformation ofSlope values for Palomar and Kitt Peak measurements to the Bus system. The vertical downward arrow is the PC2′ transformation (a constant). The magnituof these two lines is also representativeof the typical uncertainty for placement of any object within principal component space.


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Assignment of a feature-based taxonomic classificato each object was made following the description ofBusand Binzel (2002b)as originally developed byBus (1999).All taxonomic assignments are presented inTable 1. Severalof the NEOs observed in SMASSI(Xu et al., 1995)havetheir taxonomic types updated to the Bus system, but almain in the same “complex.” (For example 2074 Shoemais revised from “S” to “Sa” but remains in the S-compleFor some objects, only SMASSIR data are available othe range 0.9- to 1.65-µm. Since the Bus taxonomic claare defined over visible wavelengths below this range,not formally possible to place these near-infrared only msurements into an existing taxonomic system. Howeverspectral characteristics over the 0.9- to 1.65-µm rangegenerally recognizable as being consistent with S, C, oclasses. WithinTable 1, we list these SMASSIR-data-onresults as S:, C:, X:, where the colon denotes the uncertaof the taxonomic assignment.

Taxonomic assignments have natural ambiguity nclass boundaries, particularly for cases of low signalnoise ratio (SNR) spectra. Cases of completely ambiguand noisy spectra have been deleted from the data set. Tcases comprise only about 1% of the total presented hThe boundary between the C- and X-complexes givesto the greatest natural ambiguity for both high quality aless than perfect data. Similarly, some ambiguity can ebetween the X- and S-complexes based on the quality ospectral data for revealing the presence or absence of asorption band beyond 0.8-µm, a characteristic reflecteprincipal component PC2′. To resolve potentially ambiguous cases, our taxonomic assignments have been madeboth the principal component scores and the best matcthe spectra to the defined ranges for each class provideBus and Binzel (2002b). Asteroid 2100 Ra-Shalom providean example where high SNR data can prove ambiguouwell. The principal component scores and spectral chateristics of Ra-Shalom place it at the boundary betweenand C-types, whereBus and Binzel (2002b)denote it as Xc.With the addition of near-infrared data, the continuedcreasing slope is more characteristic of C-type than X-tobjects. For this tabulation, we place Ra-Shalom in thecategory.

Five objects have sufficiently unusual or relatively loSNR spectra that place them outside the range of theclasses or make their taxonomic assignment fully ambous. For one of these, 3908 Nyx, we follow the analysisBurbine (2000)and place it in the “V” class. For the othefour we choose to maintain the designation “U” as originalisted byBus and Binzel (2002b) and Binzel et al. (2001.(5646) 1990 TR is ambiguous between Q and V. (7888) 1UC may be an extreme form of an A-type, but falls well oside the range for this class. (8566) 1996 EN appears tan extreme form of V-type, but falls well outside the ranfor this class. (20043) 1993 EM is presented inBinzel et al.(2001a)and displays an ambiguous low SNR neutral sptrum.

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Objects observed on multiple nights or telescopesnoted in Appendix A) have their final results based oa weighted average of all measurements. For three cdatasets having significantly higher signal-to-noise (SNare preferentially used for the tabulation and analysissented here. 4660 Nereus is recognized to be an Xe-based on high SNR favorable apparition measurementstained at Kitt Peak and Palomar, as discussed inBinzel etal. (2004a). Keck measurements for Nereus show a sptral slope consistent with this classification, but have a loSNR from a faint (V20.5) apparition. For (5587) 1990 SKPNO 4-m spectra reveal an Sq type. Lower SNR SMASmeasurements are consistent with this result. For 1994 Alower SNR SMASSII data are consistent with both Sa anwith Sa tabulated inBus and Binzel (2002b). Higher SNRdata subsequently obtained at Palomar are more conswith an L-type classification that we consider a more secresult, which we tabulate here. As an additional note,Mars crosser 1011 Laodamia is reported to be an Sr-typBus and Binzel (2002b). A higher SNR spectrum obtaineusing the KPNO 4m is more consistent with an S-type dignation, which we tabulate here.

3.2. Observed taxonomic distribution

Figure 2reveals that the observed taxonomic distributof the NEO population effectively spans the full breadththe main-belt diversity, with 25 of the 26 Bus taxonomclasses represented withinthe SMASS NEO sample. Onlthe Bus class Cgh, typically found among outer main-beltteroids (beyond 2.7 AU) is not recognized within the currSMASS NEO sample. A summary of the taxonomic clasfication statistics is presented inTable 2.

Within Table 2, we also seek to collate results among mjor groups that are consistent with what are referred to“complexes” byBus (1999). We make these “complex” assignments for the purpose of establishing taxonomy inparameters to the bias-corrected population analysisformed byStuart (2003) and Stuart and Binzel (2004). Wediffer slightly from theBus (1999)use of the term “complex” in two ways. The first is our grouping the “Sq” an“Q” designations into a complex we denote as “Q.” Wmake this grouping because these types of objects aremon in the near-Earth population but are rare or absenthe Bus (1999)main-belt sample. The second is our trement of the degeneracy of the X-class. As describedTholen (1984), the designation “X” denotes spectrally simlar objects that are best distinguished by their albedos. Fhighest to lowest albedos, the distinct classes are labas E, M, P. Two factors allow us to address this degeacy. The first is the availability of albedo data for a samof NEOs (Delbo et al., 2003). The second is the correlation emerging between the E- and Xe-classes. WithinBus (1999)system, a 0.49-µm feature that distinguishesXe-class appears fully consistent with the high albedoclass, a result exemplified by 4660 Nereus(Binzel et al.,


he

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Fig. 2. Diversity of taxonomic classes observed among the near-Earth object population. Fully 25 of the 26Bus (1999)taxonomic classes are identified in tvicinity of the Earth, suggesting a broad contribution to the population from the main-belt.

Table 2Summary of NEO taxonomic classification statistics

Bus class SMASS sample Complex SMASS sample Other dataa TotalA 1 (4) A [A] 1 (4) 1 (4)B 5 (3) C [B, C, Cb, Cg, Ch, Cgh] 23 (10) 15 (1) 38 (11C 13 (2) D [D, T] 9 (1) 9 (1)Cb 3 (2) E [E, Xe] 3 (4) 1 4 (4)Cg 1 M [M] 3 3Ch 1 (3) P [P] 4 4Cgh O [O] 6 6D 4 Q [Q, Sq] 80 (11) 19 99 (11)K 7 R [R] 1 1L 7 (2) S [S, Sa, Sk, Sl, Sr, K, L, Ld] 125 (57) 20 (6) 145 (6Ld 2 (1) V [V] 14 2 16O 6 X [X, Xc, Xk] 41 (3) 6 47 (3)Q 18 (2) U [U] 3 (1) 3 (1)R 1S 76 (40)Sa 2 (6)Sk 13 (4)Sl 6 (2)Sq 62 (9)Sr 12 (2)T 5 (1)V 14X 31 (2)Xc 6 (1)Xe 2 (4)Xk 9U 3 (1)

Numbers in parentheses are a separate tabulation for the Mars-crossing (MC) population. Brackets [] depict the taxonomic classes used for the consolidationinto each complex.

a Statistics within this column are based on additional measurements ofNEOs, representing the dedicatedwork of many observers, includingMcFaddenet al. (1985), Cruikshank et al. (1991), Hammergren (1998), Hicks (1998), Rabinowitz (1998), Whiteley (2001), Angeli and Lazzaro (2002). A tabulation ofmost of these results appears inBinzel et al. (2002), for which updates and references to the original sources are maintained athttp://earn.dlr.de/nea/.

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2004a). Albedo information(Harris and Lagerros, 2002Delbo et al., 2003)is available that allows the degenerato be resolved for three objects classified as “X” inTable 1:

5751 Zao is likely an E-type while 1999 JM8 and 20BG19 are likely P-types. Similarly, 3691 Bede listed asXc-type is distinguishable by its albedo(Delbo et al., 2003

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and is likely a P-type. For these four objects plus twoobjects we tabulate separate entries among the E, M, and“complexes” in the right-hand columns ofTable 2. An in-teresting outcome of our trying to resolve the degenerwithin the X-complex is to estimate what percentage of thobserved objects might be considered “dark” and “brigalong the lines of the analysis performed byMorbidelli etal. (2002b). Within the SMASS sample, a total of 22 objechave sufficient albedo information or diagnostic spectralformation to be classified distinctly as E, M, P, Xc, Xe,Xk. Of these, 10 out of 22 or 45% fall within the P- and Xclasses that might be considered as “dark” objects, whileother 55% might be broadly considered as being “bright

Reducing the taxonomic classifications into complealso facilitates a comparison between the SMASS resand other available published data on NEOs. (Complexefectively represent a “least common denominator” betweemultiple taxonomic systems.) Results from other sources aravailable for more than 70 objects for which no SMAobservations have been obtained. (See reference fooin the table as these additional data represent the wormany dedicated observers.) Totaled together, taxonomiformation is currently available for more than 370 NEOs anearly 100 Mars-crossers. SMASS and all other progrconcur in that the S-, Q-, C-, and X-complexes account90% of all objects. SMASS differs in having a higher Xratio than that measured by other programs. This mayply be a selection effect where relatively neutral colors fromfilter measurements are most easily branded as “C-typWithin SMASS, the full visible wavelength coverageCCD spectra may better enable an X:C distinction, althothis is still subject to the X:C ambiguities described abov

A comparison between the observed near-EarthMars-crossing populations is displayed inFig. 3. Whileroughly 2/3 of all observed NEOs fall into either the S(40%) or Q- (25%) complexes, the S-complex (65%) alodominates the Mars-crossers. The low proportion of

e

complex (10%) objects among Mars-crossers may be aeffect (Section 5), with the closer proximity of NEOs allowing their discovered and spectrally measured poption to have smaller average sizes. The greater abundof D-types among NEOs may also be a size/albedo setion effect, although the apparent preferential origin oftypes from Jupiter-family comets (Section 4) and their rapiddynamical evolution with very little time spent as Marcrossers may be an important factor.

Finally, we can compare our observed taxonomic dtributions with the modeling assumptions ofMorbidelli etal. (2002b)who assumed the observed ratio of “dark”“bright” objects to be 0.165 forH < 20. Within the totalSMASS sample (Table 2), we count 207 objects havingH <

20 in the A-, E-, M-, O-, Q-, R-, S-, V-, and U-complexes (well as Xe- and Xk-class objects) as being “bright.” Amothe C-, D-, and P-complexes (as well as the Xc-class), wcount 30 “dark” objects. The 19 objects tabulated in thecomplex are problematic. We assume that 45% (9 objeof these are “dark,” where this percentage comes from“dark/bright” assessment of the X-complex discussed abCombining, we find the SMASS observed dark/brighttio amongH < 20 objects is(30+ 9)/(207+ 10) = 0.18,a slightly higher ratio compared with the 0.165 valuesumed byMorbidelli et al. (2002b). The greatest uncertaint(assuming all observational errors are random) in our 0value is our treatment of the X-complex. For example, ifconsider 9± 3 as a reasonable uncertainty for how maX-complex objects to assign to the “dark side,” the unctainty in our ratio becomes 0.18± 0.02, a range encompasing therefore compatible with the 0.165 assumed valuMorbidelli et al. ForH � 20, the SMASS observations wcalculate a dark/bright ratio following the same proceduas above. The resulting calculation yields(8+5)/(34+7)=0.32±0.06 as the observed dark/bright ratio for objects hing H � 20. (Here we assign 5 out of 12 X-complex objeto the “dark side” and assume the uncertainty in this num

ts faljectsre

Fig. 3. Comparison of the observed NEO and Mars-crossing populations, grouped into broad taxonomic complexes. Nearly 90% of the observed objeclwithin the broad S-, Q-, X-, and C-complexes. S and Q dominate the NEOs, with40 and 25%, respectively, while S dominates (65%) the MCs. V-type obrepresent the next largest category (4%) of observed NEOs, but appear absent among MCs. Relatively rare A- and E-complex objects are shown to be mocommon among Mars-crossers, each comprising 4% of the observed MC population.


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Fig. 4. Orbital element distribution of taxonomic types for near-Earth and Mars-crossing objects (shown in red). Semi-major axis (a) is in AU, inclination (i)is in degrees. Panels A and B denote main-belt asteroids as points, with curves denoting the eccentricity (e) limits for each population. Special symbols (e.∗, #, &, +) are tagged to select objects for visually correlating them from eccentricity to inclination space. TheC∗ symbols show a grouping of C-type objecneara, e, i = 2.76,0.42,29, which may be a family of Mar-crossing objects. They derive from arelatively sparsely populated high inclination region ofcentral main-belt. Panels C and D focus on rare taxonomic types that may be escapees from well known main-belt regions having corresponding rare types.For the main belt, Hungaria region asteroids are denoted by points; the denser Flora region is denoted by Floraitself; black letters denote measured types.The Mars crossers with A# and Sa# symbols (near 2.2 AU) may be olivine-rich objects that have diffused from the Flora region (containing a concentrnof similarly rare types). If they come from the Flora region, their diffusion follows the route predicted byMorbidelli and Nesvorny (1999)in which slowchanges in eccentricity occur with little or no initial change in semi-major axis. Similarly, E-type objects (labeled E& near 1.9 AU) may have follod thesame predicted diffusion paths from the Hungaria region. (The Hungaria region is rich in E-types which may be related to enstatite achondrite meteortesGaffey et al., 1992.) All tagged groupings are summarized inTable 3.

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to be 5±2.) Given the small number statistics and the untainties involved in making “dark” or “bright” assignmenfrom taxonomy alone, we believe it is an open questionto whether there is a distinct increase in the number of dobjects forH � 20.

4. Source region analysis

4.1. Orbital distribution

While transfer from the main belt to the vicinity of thEarth is a chaotic process(Wisdom, 1985), we investigatewhether the combination of taxonomic classes plus orbelements retains any signatures for tracing the originNEOs. In particular, the slow diffusion of objects into Macrossing orbits(Morbidelli and Nesvorny, 1999)may pro-

vide an observable trace.Figure 4shows the NEO and MCpopulations in semi-major axis (a), eccentricity (e), and in-clination (i) space. The alphabet soup for the NEO polation seen inFigs. 4A and 4Battests to their dynamicamixing. However these panels do show some distinctnatures remaining within the MC population, possibly csistent with slow diffusion. InTable 3we summarize oufindings for three possible groupings and discuss each gin turn below.

Figures 4A and 4Bshow a cluster of five C and C-subtyasteroids (all denoted byC∗) located neara, e, i = 2.76,0.42, 29. The dynamical evolution of these objects is likrelated to the adjacent 5:2 resonance. Whether or notgrouping constitutes a “family” (implying a collisional ori-gin from a common parent body) remains an open quesas these are common main-belt classes in this region. Whatmore, they are more diverse (comprised by C-, Ch-, an


ype

Table 3Mars-crossing asteroid groups

Group a e i Members Types Potential source

1 1.94 0.06 22 2035 Stearns, 6249 Jennifer E Hungaria region2 2.20 0.25 5 2423 Ibarruri, 3858 Dorchester,

3920 Aubignan, 5275 ZdislavaA, Sa Flora region

3 2.76 0.42 29 3581 Alvarez, 6500 Kodaira,5349 Paulharris, 5585 Parks,5870 Baltimore

C and C-subtypes No identified source, but a common tfor outer main belt

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types) than is typical for a homogeneous family. Unlikeother groupings inTable 3, this one has no apparent assoation with any previously known cluster.

Figures 4C and 4Dfocus on the inner belt and taxonomclasses (A, E, Sa, V) that are generally interpreted to indiobjects that have undergone a moderate degree of heand differentiation. Being rare taxonomic types, they othe opportunity for some traceability in their dynamical evolution. According to the diffusion model ofMorbidelli andNesvorny (1999), evolution in eccentricity occurs first withlittle or no change in semi-major axis until Mars-crossorbits (with the potential forMars encounters) are reacheWe find two cases of rare taxonomic types that may be trable across the Mars-crossing boundary as direct evidenthis predicted diffusion. In the first case, we find that E-tyasteroids (often related toenstatite achondrite meteoriteGaffey et al., 1992) are abundant in the Hungaria regilocated neara, e, i = 1.94,0.06,22. Two E-type objects apparently caught in the act of diffusing from the Hungarias2035 Stearns and 6249 Jennifer (denoted by the E& symin Figs. 4C and 4D). A third object 2449 Kenos (denoted bE+) matches in eccentricity, but not inclination. If relateit has already begun a direct interaction with Mars affectits orbit plane.Gaffey et al. (1992)make a strong case foE-type 3103 Eger to be derived from the Hungaria regwhere its evolved orbital elements(a, e, i = 1.40,0.35,20)place it completely out of the area ofFigs. 4C and 4D, requir-ing a much more evolved dynamical history than the objewe list in Table 3.

A second case for diffusion in orbital eccentricity to bcome Mars-crossing objects is also found in the A andrich Flora region neara, e, i = 2.2,0.16,6. The visible spectra of A- and Sa-type objects suggest a mineralogy ricolivine, as does the Sr-class also seen in this region(Gaffey,1984; Florczak et al., 1998)A group of four objects residjust across the Mars-crossing boundary near 2.2 AU. Infigure they are denoted by Sa# and A# and their identare given inTable 3. All four have the same inclinations aFlora, making them prime candidates for having recently diffused from this region but not having yet begun substaninteractions with Mars. Taken together, the apparent traof diffusion from the Hungaria region and the Flora regprovide strong observational support for the weak resonadiffusion into the Mars-crossing population predictedMorbidelli and Nesvorny (1999).

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Figure 4(as well asFig. 2 and Table 2) indicates thatV-types are rare or absent among observed Mars-cros(16 observed V-types fall in the NEO category, none areserved in this sample among MCs.) This dichotomy impa low eccentricity main-belt origin for the V-types via theν6 and 3:1 resonances, a process in which objects passrapidly (and therefore unlikely to be observed) throughMars-crossing phase in becoming NEOs(Morbidelli et al.,2002a). High eccentricity objects can become NEOs viaslow diffusion process ofMorbidelli and Nesvorny (1999),with a long residence time as Mars crossers. The raritV-types among Mars crossers implies they do not havsignificant high eccentricity source. Most scenarios (e.gConsolmagno and Drake, 1977; Binzel and Xu, 1993) pre-dict that V-type NEOs (and HED meteorites) are derivfrom Vesta or its apparently associated collision fragmeThe low orbital eccentricities of Vesta and the “Vestoids,indeed they are the source of V-type NEOs, would prodthe predominance of V-type NEOs relative to Mars crossthat is observed.

Finally we note a distinct variation in taxonomic distbution with respect to the jovian Tisserand parameter(T ),as readily apparent inFig. 5. We describe this more fully inSection 4.3.

4.2. Taxonomic signatures from source regions

With the available data set of measured taxonomic perties for the NEO population, a marriage of observatiwith theoretical modeling now becomes possible for atailed examination of NEO sources.Bottke et al. (2002)provide a model for calculating the probability of an oject entering NEO space from one of five source regiotheν6 resonance, the Mars-crossing zone, the 3:1 resonathe outer belt, and Jupiter family comets. For each objecour sample, Bottke (personal communication) providedmodel values for the five source regions based on the objcurrent orbital elements, where the sum of the five moprobabilities is equal to one. We couple these model probbilities with the NEO data consolidated into the “complexeof Table 2. We examine the resulting source region probility distribution for each complex separately. For exampwe take the Bottke model probability numbers for just thecomplex objects and sum the fractional probabilities witeach of the five source bins. The results for each of thfive source bins are normalized so that their total proba


le).t the

Fig. 5. Distinct differences are found between theT � 3 andT > 3 populations when comparing their fractional abundances (number in class/total sampC-, D-, and X-complex objects (Table 2) are found to dominate forT � 3, all of which are in low albedo categories (assuming the “X” objects represenP-class), consistent with the findings ofFernandez et al. (2001). TheseT � 3 objects may be candidates for extinct comets(Wiessman et al., 1989, 2002).

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ity is one. We follow this in turn for D-, E-, M-, P-, Q-, SV-, and X-complexes, whereFig. 6A displays the results foeach complex. For the objects which we sampled, theν6 res-onance prevails as the source with the greatest contribuaccounting for 46% of the source probability for the sumour entire sample (all complexes combined).Bottke et al.(2002)(see their Table 3) predict a steady state NEO ctribution of about 37% from theν6 resonance. We accoufor the greaterν6 dependence in our sample as being duobservational selection effects: objects at the inner edgthe asteroid belt are more easily discovered and more lito have their properties measured in a magnitude limsurvey. Source probabilities for our sample from the inmediate Mars-crossers (27%), the 3:1 resonance (19%)the outer belt (6%) match well with the values ofBottkeet al. (2002). Our sample has a lower (2% as compared6%) source probability from Jupiter family comets, an efflikely due to the difficulty of discovery and observationhigh eccentricity and typically distant objects, especiallthey are dominated by classes having low albedos.

While differences in the source regions for one compcompared to another are apparent inFig. 6A, they are morereadily seen inFig. 6Bwhere we normalize each sourcegion by the average for the total sample. As an examC-type objects contributed by the 3:1 resonance have atogram value 0.25 inFig. 6A, while the average contributiofrom the 3:1 resonance for our entire sample is 19%.Fig-ure 6Bshows the normalized C-type contribution via theresonance to be given as 1.31 (leftmost striped bar), w1.31 is the quotient of 0.25 normalized (divided) by 0.If all taxonomic types were equally likely to be contributfrom all sources, all histogram values ofFig. 6B would beunity.

Distinct variations from unity are revealed inFig. 6Bthat represent the signatures of higher (or lower) thanerage contributions to the NEO population from the fi

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d

-

sources. The greater than unity values for the 3:1 and obelt sources within the C-complex reveal that these regdeliver a proportionally larger fraction of C-type objectsnear-Earth space. Such a finding is fully consistent withpredominance of C-type objects at the 3:1 resonance anyond. Similarly, the dominance D-types being contribufrom Jupiter family comets is strikingly shown, with impotant implications for deriving the extinct comet fraction frothe NEO population, as discussed below. Because S-typcomprise the greatest fraction of our total sample, it can bexpected that their source contributions will closely appromate the average. The high degree of similarity seen betwthe S- and Q-complexes may have an important implicatif S- and Q-asteroids are related (such as Q-types just bein“fresh” S-type surfaces), then a necessary (but not sufficcondition is that they show similar structures for their souregions. V-types also show a very similar profile relativethe S-types. While theFig. 6Bhistogram predicts relativelequal entry of V-type NEOs via the Mars-crossing routeresonance routes, the lack of observed V-type Mars cro(noted above) suggests a quick passage from a resonangin through the Mars-crossing phase, into NEO space.

Finally, in our examination of source region signatuwe evaluate the E-, M-, P- and X-complexes. X-complexjects show a very strong contribution to the JFC populathat we interpret as arising from the fraction of low albeP-types, which by definition are part of the X-complex. Wnote that only a total of 10 X-type NEOs have albedoformation for discriminating them into the E, M, P class(Table 2). Thus the E, M, P signatures inFig. 6B are basedon a very limited sample and the resulting implicationsonly be considered preliminary. Among E-types, the higthan average contribution from the innermostν6 resonanceis consistent with the high abundance of E-types in thejacent Hungaria region. Both M-and P-types show evidencfor high concentrations from the outer belt, where primit


6,is

hesedashed

senance

es

Fig. 6. A. Contributions by taxonomic complexes (Table 2) to the NEO population are presented as a function of the source regions modeled byBottke et al.(2002). These regions:ν6 resonance, Mars-crossing (MC) zone, 3:1 resonance, outer belt (OB), and Jupiter-familycomets (JFC), respectively, contribute 427, 19, 6, and 2% of the sampled population, where these percentages are depicted by the symbols on the left. For each complex, the sum of all sourcesunity. Theν6 resonance is seen as the most prolific provider for all classes. B. Signatures of source region contributions for each taxonomic complex. Tsignatures are revealed by rationing (e.g., by 0.46 forν6) the contribution in each zone by its average and comparing the results with respect to unity (line). Many distinctive sources are revealed: C-types have a proportionally higher contribution from the 3:1 resonance and outer belt while D-typesare stronglysourced from Jupiter family comets. The sample size for E-, M-, P-types is small, thus their trends are only preliminary indications: E-types preferentiallyenter via theν6 resonance. M-types in the sample show an originpreference in the outer belt, perhaps suggesting they may be examples of “primitive” objectwithin this class(Rivkin et al., 2002), or that they are capable of long lifetimes and substantial orbital evolution if metallic. P-types also show a provfrom the outer belt. X-types (which by definition contain an unknown fraction of low-albedo P-types) show a high relative contribution from the outer belt.Because they dominate the total sample size, S-class objects can be expectedto reflect the average. Both Q- and V-types also reflect the average, suggting
similar main-belt region origins as S.
pesypeple

ter-udealueto

ectsonalwithiter.

ashe

(not strongly heated) asteroids predominate. While P-tyare known to predominate in the outer belt, seeing an M-tspike in the outer belt is confusing, but is based on a samof only three objects. We discuss the M-types inSection 6.

4.3. Tisserand parameter and comet fraction

In a classic analysis,Fernandez et al. (2001)found a de-pendence on albedos related to the Tisserand parameterT

(with respect to Jupiter):

T = aj

a+ 2 cos(i)

√(1− e2

) a

a,

j

whereaj is the semi-major axis of Jupiter. Inner and ouSolar System objects havingT � 3 were found to have distinctly lower albedos, which led Fernandez et al. to conclthese objects were candidates for extinct comets. The vof T � 3 denotes objects that are dynamically coupledJupiter, as is the case for Jupiter family comets. Objmay change their Tisserand value through non-gravitatiforces (such as cometary outgassing), by interactingother planets, and as a result of the eccentricity of JupBecause NEOs do interact with inner planets,T � 3 is nota rigid boundary for objects that may have originatedJupiter family comets. Similarly, objects originating in t


hatrig-

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main-belt withT > 3 can undergo planetary encounters tcast them into the control of Jupiter. Thus objects that oinated as Jupiter family comets may now haveT > 3, justas objects originating in the main belt can now haveT � 3.Cometary outcasts across this border may be recognizesuspected) based on their very low albedos or unusualnomic classes, as discussed below. Asteroidal outcasts mmost likely manifest themselves inT � 3 orbits as S- and Qtype objects—as these most common types of NEOs whave the greatest random chance of being scattered atheT = 3 boundary.

Here we make a similar but independent analysis toof Fernandez et al. (2001), utilizing T � 3 as a distinguishing parameter but focusing only on the near-Earth polation. Rather than albedo, we use taxonomic classeour variable. Current discovery statistics (IAU Minor PlanCenter) show that 7% of all catalogued NEOs haveT � 3.Within the sample population of SMASS, 6% of the oserved objects haveT � 3, a close representation for the dcovered population. While C-, D-, Q-, S-, and X-complehave more than oneT � 3 member in our sample, ourT � 3sample of NEOs is dominated (50%) by C-, D-, and X-clobjects as shown inFig. 5. For our analysis of theT � 3population, we make two assumptions: First, theT � 3 S-and Q-types within our sample are randomly scattered mbelt objects as discussed above (where S and Q are thelikely letters to be drawn from the main-beltScrabble® bag).Second, we assume the X-type objects havingT � 3 are ac-tually P-types based on their very strong signatures of bderived from the outer belt (Fig. 6B). (Of the 6 X-complexNEOs within ourT � 3 sample, we have albedo informtion for only one, 1999 JM8. From radar measurementBenner et al. (2002), 1999 JM8 is inferred to have a lovisible wavelength albedo.) This predominance of the Cand (assumed) P objects forT � 3 is consistent with thelow albedo correlation shown byFernandez et al. (2001.Starting with 7% of the discovered NEO population hing T � 3, and our finding of 50% of these being likelow albedo classes, we estimate that 4% of the totalcovered NEO population has low albedos based on tfactors alone. This is lower than the 9% value estimatedFernandez et al. (2001), although their sample considerboth NEOs and objects in unusual (but not near-Earthbits.

The observed population of NEOs having bothT � 3 andmeasured (or inferred from taxonomic class) low albedosconsidered prime candidates for extinct comets, as revieby Weissman et al. (1989, 2002). Deriving the actual population of objects having bothT � 3 and spectra presumabcompatible with extinct comet origins must take intocount the discovery bias against these objects arisingtheir inferred low albedos.Stuart (2003)(updated byStuartand Binzel, 2004) performs this diameter limited bias anasis and finds that 30± 4% of all NEOs haveT � 3, where85± 15% of these have C-, D- or P-type taxonomies. Thtwo factors (0.30× 0.85) yield an estimate of 25± 5% for

r-t

s

st

the total proportion of NEOs that have bothT � 3 and C-,D-, and P-type spectroscopic properties above any giveameter. Limiting our attention to just D-types, for which tdiameter limited bias analysis ofStuart (2003)finds themto be 43± 19% of all T � 3 objects, yields an estimate13± 6% (equal to 0.30× 0.43) for the total proportion oNEOs that have bothT � 3 and D-type spectroscopic proerties. We use this range of 13–25% for the percentagNEOs above any given size having both orbital and sptral characteristics being compatible with extinct comets aa starting point for our discussion inSection 6.

5. Size dependence of spectral properties: evidence forspace weathering

While NEOs generally fall within regions of principcomponent space (Fig. 1) populated by main-belt examples, many NEOs are uniquely found to populate thetype region and the intervening spectral component sbetween the S- and Q-types(Binzel et al., 1996). Q-typeshave long been associated as spectral analogs to ordchondrite meteorites, with 1862 Apollo being the protype example for this class(McFadden et al., 1985). Thepossible link between the most common asteroids (S-tyand the most common meteorites (ordinary chondrites)been long debated (e.g.,Wetherill and Chapman, 198Chapman, 1996). Thus, revealing a connection betweenS- and Q-type asteroids has the potential to establish thesought link between S-type asteroids and ordinary chonmeteorites.

To analyze this possible connection, we utilize the sset of near-Earth and Mars crossing objects having Q-,and S-types as determined using principal componentssured within the SMASSII system(Bus, 1999)as compiledhere in Table 1. We restrict ourselves to consider thetypes, Sq-types, and the S-type core of the S-complexThe component showing the greatest difference betweeQ-, Sq-, and S-types is the Slope parameter, defined aaverage slope of the spectrum over the wavelength inval 0.44- to 0.92-µm(Bus, 1999). To explore a possiblsize dependence, we convert the H magnitudes inTable 1to diameters by utilizing the albedos determined throthe thermal modeling of this same population(Delbo et al.,2003). For objects with no specific albedo determination,use the mean albedo(pv) values (0.244 for S-types; 0.25for Sq- and Q-types) from theDelbo et al. (2003)sample toestimate diameters from:

D = 1329(pv)−1/210−0.2H .

While both random and systematic errors may be prein the parameters used to estimate diameters, these eshould have little effect in the broad statistical analysisapply here.

The data points inFig. 7 display the diversity of spectral slopes versus diameters for Q-, Sq-, and S-types, w


for

e

ld

Fig. 7. Data points (open circles) show measures of spectral slopes (determined over 0.44- to 0.92-µm) versus diameter for near-Earth and Mars-crossingobjects residing within the S-, Sq-, and Q-classes ofBus (1999). A running box mean is shown by filled squares (box size= 50, with error bars depictingthe standard deviation of the mean). The running box trend asymptotically approaches the mean slope (dashed line)for SMASSII main-belt S-type asteroids,reaching this limit at a size of 5 km. It appears that 5 km may represent a “critical size” in the evolution from ordinary chondrite-like (Q-type) to S-typesurfaces, depicted by the Q→ S vector. (This vector corresponds to the Slope difference between the Q-type 1862 Apollo and the main-belt averageS-types.) The vectors labeled “H,” “L,” and “LL” show the effects of a reddening model for ordinary chondrite meteorites resulting from the addition of 0.05%submicroscopic iron (SMFe). The magnitude of the transition for the meteorites is comparable to the magnitude of the Q→ S vector for the asteroids. ThSMFe model vectors were determined by calculating the slopes of meteorite spectra before and after applying the 0.05% SMFe curve fromPieters et al. (2000),using the same method asBinzel et al. (2001b). All meteorite spectra are fromGaffey (1976)where “H” is the average for H6 chondrites, “L” is the BaMountain L4 chondrite, and “LL” is the Olivenza LL5 chondrite.

cat-lestwerlessboxob-

eantan-singenceoted

ula-toan

eterds

oids

”as-

anallyper-widen obthern the

stu-hin, noag-

as a

ayenernd

e aticion:ass ofErosel-g-t al.

asedult is

, as

pacepic

with

the errors discussed above simply contribute to their ster. Most evident is the higher dispersion for the smalobjects, although we note that at larger sizes their fenumbers would tend to make any comparable dispersionapparent. To search for trends, we employ a runningmean (box size 50) stepping through the sample oneject at a time from smallest to largest. The resulting mvalues (with the overall sample variance reduced to the sdard deviation of the mean) show a clear trend for decreaspectral slope with increasing size. While a size dependin spectral properties for S-types has been previously n(e.g.,Gaffey et al., 1993a, 1993b; Rabinowitz, 1998), thevery small diameters sampled herein for the NEO poption reveal a new characteristic: over the range of 0.15 km, mean spectral slopes appear to increase sharplythen asymptotically approach a value of 0.44 near a diamof 5 km. Most interestingly, this value of 0.44 corresponto the mean spectral slope of all main-belt S-type asterwithin the SMASSII sample ofBus (1999). One possible in-terpretation ofFig. 7is that 5 km represents a “critical limitat which a size dependent transition for Q-type to S-typeteroids reaches “completeness.”

We first evaluate whether this trend is real or due toobservational selection effect. Smaller objects must typicbe closer to the Earth in order to have their physical proties measured. Being closer, these objects can have arange of phase angles (Earth–asteroid–Sun angles) thajects located farther away. An analysis to determine whephase angle induces any effect (such as reddening) o

d

r-

spectral slope was conducted by MIT undergraduatedent Nancy Hsia. This analysis found no correlation witour sample between phase angle and Slope. Similarlycorrelation was found between Slope and the observed mnitude, thus revealing no systematic effect in our resultsfunction of limiting magnitude for the observations.

We consider two possibilities for real effects that mcause the trend inFig. 7. The first is related to surfacparticle sizes. Nominally, larger bodies should have firegolith characteristics owing to their greater gravity alonger surface evolution lifetimes. (Larger bodies havlonger lifetime in between collisions sufficiently energeto catastrophically disrupt them.) Thus we ask the questcan the trend inFig. 7 be due to decreasing particle sizepart of the process of regolith development? An analysialbedo and spectral contrasts within Psyche crater onby Clark et al. (2001)provides a case against regolith devopment causing the increasing spectral slope with increasindiameter. For an olivine, orthopyroxene, plagioclase mixture intended to model Eros (an S-type asteroid), Clark e(see their Fig. 18) found that finer grain sizes had decrespectral slopes relative to coarse grains. This model resopposite to the trend withinFig. 7.

The second possibility is that the age of the surfaceaffected by “space weathering” (see review byClark et al.,2002) causes the observed trend. Current models for sweathering being due to the deposition of submicroscoiron (SMFe)(Hapke et al., 1975; Pieters et al., 2000)sug-gest that older surfaces become increasingly reddened


ds

Fig. 8. Albedo versus diameter for S-, Sq-,and Q-class near-Earth and Mars-crossing objects. Data are from the tabulations byBinzel et al. (2002), Harris anLagerros (2002), and Delbo et al. (2003). Individual values are plotted by open circles with a running box mean based onn = 5 for the box size. (Error bardepict the standard deviation of the mean.) The dashed line depicts the average albedo for S-class asteroids in the main-belt(Harris and Lagerros, 2002).

H,erestenids.rgerhavemasto

pery ares

er(in-lye efringcalemat-roid

ocesofids.ce,

withd by-

ingible

etsas

ys

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lasstors

hinery

ntlymicn ofnesical,ly-R

ere.of

st

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the increasing accumulation of SMFe over time. Vectorswithin Fig. 7 show the effect on the Slope parameter forL, and LL chondrites by adding just 0.05% SMFe, whthe direction and magnitudes of these vectors are consiwith the apparent transition from Q-type to S-type asteroSmall asteroids have shorter collisional lifetimes than laones, suggesting that on average smaller asteroids willyounger, fresher, and less reddened surfaces. Not all sasteroids must have fresh surfaces since collisions are achastic process. However, the dispersion in surface proties may be expected to be higher at small sizes since themost likely to display “fresh” surfaces in addition to surfacthat are stochastically “old.” Figure 7indeed shows a greatdispersion at the smaller sizes and a reddening trendcreasing spectral slope) with increasing size and presumabolder average surface age. If this trend is indeed an agfect, the most important implication for the space weathehypothesis is that the apparent limit at 5 km sets a timesover which the responsible process effectively reachesuration. In other words, the average age of a 5 km astesurface sets the time scale for the space weathering prto be effective in “transforming” the spectral signaturesordinary chondrite like bodies into those of S-class asteroFigure 8suggests, but does not offer convincing evidenthat 5 km is also a transition size for decreasing albedoincreasing size. Revelation of an albedo trend, discusseDelbo et al. (2003), remains difficult due to the limited number of such measurements for NEOs.

6. Discussion

In this section we expand on our findings for constrainthe extinct comet fraction, NEO source regions, and possspace weathering trends.

t

ll--e

-

s

Estimates for the percentage of extinct or dormant comwithin the NEO population have ranged from as high∼ 50%(Wetherill, 1988)to the currently predicted vicinitof 2–10%(Bottke et al., 2002), where the latter estimate wamade for a specific magnitude range, 13< H < 22. As de-scribed byWiessman et al. (2002)(which incorporates thBottke et al. (2002)findings), physical studies combinewith dynamical parameters have provided key informafor constraining this percentage and for identifying indivual objects as specific extinct comet candidates. Apart fthe goal of simply understanding the make-up of the Npopulation, identifying extinct comet candidates and demining their overall fraction provides an understandingthe end states of comets.

Two quantities, Tisserand parameter and taxonomic c(and/or measured albedo) are currently the best indicafor recognizing bodies that may be extinct comets witthe NEO population. Current magnitude limited discovstatistics show only 7% of all NEOs haveT � 3, and inSec-tion 4we found roughly one-half of these (and conseque4% of the total observed NEO population) have taxonoproperties consistent with low albedos. The combinatioeccentric orbits forT � 3 objects and low albedos combias strong bias factors against their discovery and phycharacterization.Stuart (2003)(updated inStuart and Binzel2004) performs a diameter limited bias correction anasis utilizing extensive NEO search statistics from LINEA(Stokes et al., 2002; Stuart, 2001), NEO albedos(Delbo etal., 2003), and the taxonomic distributions presented hThe resulting diameter limited estimate for the fractionthe NEO population residing inT � 3 orbits is 0.30± 0.04.Diameter limited bias-corrected fractions for theT � 3 pop-ulation having low albedos within C-, D-, or P-type or juD-type taxonomies are 0.85± 0.15 and 0.43± 0.19, respec-tively. All of these factors are summarized within columa–f of Table 4, where high and low values for these pa


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meters are from the one-sigma uncertainties arising fPoisson sampling statistics or from the formal error analywithin the models. For all calculations progressing throuthe columns ofTable 4, these errors are propagated assuing they are independent and normally distributed.

There are a variety of ways to examine the factors witTable 4in order to achieve a new estimate for the exticomet fraction based on current discovery, albedo, andonomy statistics. Allowing that C-, D-, and P-type objeare compatible in terms of spectra and inferred albedoextinct comet nuclei, then an estimated 25% (determifrom 0.30× 0.85) of the NEO population above any givediameter has properties consistent with being derived fextinct comets. Restricting the extinct comet candidatethose having the lowest inferred albedos (D-types), yield13% estimate (0.30× 0.43) for candidate extinct cometsthe population. Factoring in the associated uncertaintiesTa-ble 4, columns g and h) expands the range for both estimconsiderably.

Having aT � 3 orbit and taxonomic properties constent with a low albedo, however, is not sufficient to implyobject is an extinct comet. One of the complicating factorsis that low albedoasteroidscan be scattered intoT � 3 or-bits, thereby contaminating the sample of candidate excomets. From the model ofMorbidelli et al. (2002b)(alsoA. Morbidelli, W. Bottke, personal communication, 2003for any given diameter or larger, 35% of the low albedojects within theT � 3 population are expected to be derivfrom asteroids scattered mostly from the outer belt. Thuslow albedoT � 3 objects, the fraction being extinct comeis estimated to be 0.65± 0.10. Correcting theT � 3 candi-date comet fraction (Table 4, columns g and h) by the 0.6factor yields estimates of 8 to 16% for the diameter limited extinct comet fraction within the entire NEO populatioHowever, in reaching a final estimate, the fraction of the extinct comet population scattered fromT � 3 toT > 3 cannotbe ignored.Wiessman et al. (2002)(Table 1 and Table 3list many extinct comet candidates (such as 3200 Phaethaving Tisserand values greater than 3. We estimate thetion of the extinct comet population scattered toT > 3 to be0.02± 0.02, where this is an ad hoc value intended to recnize the potential existence of these scattered objects wallowing that their numbers may be small. With this finparameter we reach the two rightmost columns ofTable 4giving 10–18% as our best estimate for the NEO poption, for any given diameter or larger, being comprisedextinct comets. Propagating our uncertainties, formallyestimate is 10± 5% if only D-type objects are consideredcomet candidates and 18± 5% considering all dark classe(C, D, P).

How does our 10–18% range compare with the previestimates? Direct comparisons must be done carefullycause of different debias criteria. Our results are diametelimited, meaning above any given size, 10–18% of the Npopulation is estimated to be extinct comets. The 2–10%sult of Bottke et al. (2002)is modeled over the magnitud


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range, 13< H < 22. Accounting for the differences betwemagnitude limited and diameter limited bias corrections (discussion byMorbidelli et al., 2002a) increases the ranggiven by Bottke et al. These authors (A. Morbidelli, W. Btke, personal communication, 2003) find their diameter lited extinct comet contribution to be 17%, a value quite csistent with the estimate we report here.

While a compositional gradient has been well knowithin the asteroid belt(Gradie and Tedesco, 1982)and avariety of asteroidal source regions have been identifiedthe NEO population (e.g.,Wetherill and Chapman, 1988Greenberg and Nolan, 1989; Morbidelli et al., 2002a), theirchaotic routes to the inner Solar System(Wisdom, 1985)may give little expectation of recognizable signatures liing NEOs to their sources. With the contributions ofSMASS program and many other observers, the sampleof spectral and taxonomic properties of the near-EarthMars-crossing populations has now grown large enougbe investigated for distinct signatures. Key to this invegation is the statistical model ofBottke et al. (2002)thatassigns probabilities for an object’s source based on itsrent orbital elements. The source region signatures revewithin Fig. 6, with one exception (M-types, discussed blow), show results consistent with the basics of theGradieand Tedesco (1982)compositional gradient within the ateroid belt. This consistency gives a mutual check betwthe observational statistics from SMASS and other obserand the source models ofBottke et al. (2002). E-types,known to dominate the inner belt Hungaria region, shopredominant inner belt source region signature, and arect “tracer” may be most apparent for objects just acrossMars-crossing boundary. (Table 3lists candidates for addtional objects that may be traced to distinct sources baseorbital and spectral characteristics.) S-types show a cleanature from the inner main-belt where they are well knoto dominate. The indistinguishable source signature ofQ-types with the S-types is a strong argument for their cmon origin, consistent with their spectral difference beingartifact of a process such as space weathering(Clark et al.,2002). C-types and P-types very nicely show their strongsource region signatures from the outer asteroid belt, anD-type signature from Jupiter family comets constrainscomet population, as discussed above.

The occurrence of V-types among NEOs but theirity among Mars-crossers serves as a dynamical trthat specifically constrains the orbital eccentricity of theisource(s). While high eccentricity objects can evolve iNEOs with a long (and likely to be observed) residence tas Mars-crossers, objects in initially low eccentricity orbspend very little time as Mars-crossers in the course of t“fast track” evolution through theν6 and 3:1 resonance(Morbidelli et al., 2002a). Thus low eccentricity sources apear the most likely source for the V-types, a condition mby Vesta and its likely association of “Vestoids”(Binzel andXu, 1993). Based on the orbital distribution of V-types

-

their own sample,Dandy et al. (2003)come to the sam“fast track” conclusion for the V-types.

As noted inSection 4, the M-types provide the most cofounding source region signature inFig. 6, suggesting a predominant origin from the outer belt. Such an origin doesfollow the colloquial “metallic” interpretation for “M,” morelikely to come from the more strongly heated region ofinner asteroid belt. However,Jones et al. (1990)as well asothers (e.g.,Rivkin et al., 1995, 2000, 2002) have found ev-idence of water of hydration in some M-types, implying a“metallic” interpretation of their composition solely bason visible wavelength data (and albedo) may be susp(The mnemonic for “M” may be best characterized byword “muddle.”) While it is tempting to simply attribute thouter belt spike for an M-type source to the revelation of “drated M-types” within the outer belt, we strongly cautagainst any premature conclusion since this spike for thetypes withinFig. 6is based on a sample of only three objecStrongly influencing this spike is the M-type object (6171986 DA, whose orbital elements(a, e, i = 2.81,0.58,4.31)lead to aBottke et al. (2002)model probability of 0.81 forcoming from an outer belt source. 1986 DA is an objectwhich the interpretation of high metal content seems secas it exhibits one of the highest radar reflectivities measto date(Ostro et al., 1991). We suggest that the outer besource probability (if meaningful, since there is still a 19chance for coming from other regions) for the origin of 19DA may be a direct consequence of this object’s strengthpossible long-term dynamical evolution within the main-bprior to being cast into near-Earth space. Assuming 1DA is indeed “metallic,” it may have had a particularly lomain-belt lifetime against collisional disruption, as corroorated by the very long cosmic-ray exposure ages ofmeteorites(Buchwald, 1975). As a counterpoint to the “mostraceable” objects being those just over the Mars-crosseeccentricity border, the least traceable to their early SolaSystem formation location may be high strength objectscan have very long survival lifetimes and experience sigcant Yarkovsky drift within the main-belt prior to enteringresonance transporting them to near-Earth space.

There is now a large body of spectral evidence linkS-type asteroids to Q-type asteroids, implying a linktween S-asterods and ordinary chondrite meteorites (Binzel et al., 1996; Chapman, 1996; Rabinowitz, 19Trombka et al., 2000). First we caution that objects denotas “S-asteroids” span a wide range of mineralogies fthose being highly pyroxene-rich and olivine poor (thesub-class ofGaffey et al., 1993a) to those being dominateby olivine (the Gaffey SVII sub-class). For this reasonSection 5we focused our search for S-type to Q-typelationships just along the transition in spectral slope frQ-types to Sq-types to the core of the S-types so as tminimally effected by variations in mineralogy. We findsize dependence to spectral properties, where Q-type aoids are more common at smaller sizes, a result previonoted by others, includingBinzel et al. (1996), Hammergre


ey

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(1998), Hicks et al. (1998), Rabinowitz (1998), Whitel(2001), and Dandy et al. (2003). Our conclusion that thisdiameter dependent trend is largely caused by “space wering” effects on spectral slope, rather than by variationsmineralogy, stands to be verified or refuted by follow-up oservations for these objects over near-infrared wavelenAt present, the consistency of the magnitude of the ovechange with the effects of a minor amount (0.05%) of subcroscopic iron (Fig. 7), adds weight to this particular modfor space weathering.

Our observations (Fig. 7) point to a key transition occurring at 5 km, where an increase in slope with increasingmatches up with the average value for main-belt S-typ(Main-belt S-types have been measured over the size rof ∼ 10 to several hundred km.) An inflection in the runing box trend occurring at 2 km may be the characteriin spectral colors noted byRabinowitz (1998)based on ananalysis of filter photometry. At present it is not clear whsignificance (if any) there is for the 2 km inflection asFig. 7shows that the dispersion of the raw data remains simover the full range from∼ 100 m to 5 km. The transition a5 km appears more robust in that there is a marked decrin dispersion of the raw data. The (perhaps now “classicinterpretation of a size dependent trend is that we are seobjects with increasing average collisional ages and thfore increasing average surface exposure times, i.e., incingly “weathered” surfaces(Gaffey et al., 1993a; Binzel eal., 1996, 1998; Rabinowitz, 1998). The observed transitioover the range of 0.1 to 5 km (Fig. 7) suggests we are seing a “completion” of the space weathering process. If 5represents the size at which the effects of space weaing become “complete,” then 5 km may represent a criticasize where the timescales for two independent processematched. The first is the timescale over which collisionscavate and refresh the surface with unweathered mateThis timescale must be less than or equal to the collisiodisruption age for 5 km objects, for which estimates rafrom 107 years(Farinella et al., 1998)to 109 years withinthe main-belt(O’Brien and Greenberg, 2003), depending onmodels for the body’s impact strength.Cheng (2004)arguesthat objects larger than 5 km are survivors over the age oSolar System while those smaller are second (or later) geation fragments that must have younger collisional agestheir surfaces. The second timescale is the interval for dsition of sufficient submicroscopic iron (as an example ospace weathering process) to alter the slopes from ordichondrite material to S-type asteroids.

The “classical” model of surface exposure time fahowever, if the effective timescale for space weatherinextremely short compared with collisional timescales.other words, if space weathering is so rapid as to bestantaneous” compared with the interval over which anteroid’s surface is refreshed, then surface exposure airrelevant. (If the youngest surfaces of the smallest boare instantly weathered, there is nothing about surfaceor collisional age that distinguishes them from the olde

-

.

e

-

-

e

.

-

and also weathered surfaces of larger bodies.) Current mels suggest that space weathering processes indeed meffective very rapidly, perhaps in as little as 50,000 ye(Hapke, 2001), a time that appears very short comparedcollisional timescales. We note that the possible very yodynamical age of the Karin family(Nesvorny et al., 2002could provide insight into space weathering timescales.

For the case of rapid space weathering, the 0.1 to 5trend inFig. 7, with an apparent “completion” at 5 km, requires an alternate explanation from that of surface agepropose that the trend is a measure of increasing regolithvelopment, driven both by the average surface age andincreasing gravity (necessary for regolith retention) ofbody. (Generally decreasing rotation rate, with increassize, may also be a factor for regolith retention.) Underscenario, space weatheringcannot commence until a regolibegins to develop. The possible inflection at 2 km noabove(Rabinowitz, 1998), could be the start of sufficient regolith development or retention and therefore the onseweathering effects. As the abundance of regolith grows,instantly weathered, and creates surface reflectance prties that increasingly change from being ordinary chondrlike to being like S-asteroids. The “completeness plateau5 km may be the size where there is sufficient surfacelution and gravity to begin to retain regolith (or a particuparticle size distribution for a regolith) in a manner thatmains consistent with the regolith properties of 10–100main-belt S-type asteroids.

7. Conclusion

The richness of our increasing scientific understandinthe near-Earth object population is a natural product ofnecessity to characterize this population toward the pracal goal of defining their size distribution and impact hazar(Stuart, 2001, 2003; Stuart and Binzel, 2004). Our scientificinterest in the NEO population is driven by our desire torive insights to their asteroidal and cometary sources anunderstand asteroid–meteorite connections. The NEO stral data set is now growing large enough to correlate phical properties with dynamical source regions. The corrtion of low albedo D-, C-, and P-types with Jupiter famcomet sources is an important synergy between dynammodels (e.g.,Bottke et al., 2002) and physical observationrevealing a strong source signature for comets withinNEO population. Taking into account the best availablecovery, taxonomy, albedo, and bias correction models,estimate that 10–18% of the NEO population (at or abany given diameter) may be extinct comets.

With the growing spectral data set and source region mels, correlations between S-type asteroids and Q-type (nary chondrite-like) bodies continue to build. Both clashave identical asteroid source region profiles. A clear sdependent transition appears to forge a link betweenasteroids and ordinary chondrites, where surface expo


h al-eath. If

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age, the process of regolith development, and/or regolitteration (perhaps by submicroscopic Fe as the space wering agent) completes the transition at sizes of 5 kmthis interpretation is correct, a substantial proportion otype asteroids 5 km and above may be composed ofnary chondrite-like materials, but have mature (reddensurfaces giving rise to spectral mismatches with ordinchondrites. In situ measurements by the NEAR-Shoemspacecraft of the S-asteroid Eros(Trombka et al., 2000)givethe most direct and independent evidence for this link todinary chondrites.

Continued reconnaissance of the NEO population willveal additional unusual compositions, raise new questionfor the detailed understanding of their origins, uncovermerous additional extinct comet candidates, and providether direct evidence for the spectral evolution of asteroisurfaces that is necessary to fully unravel remaining qtions about asteroid–meteorite connections. Having a wsampled population enables a scientific basis for choofuture spacecraft mission targets among these relativelyily accessible worlds. Perhaps most pragmatically, all thawe learn scientifically is of direct practical benefit to the uderstanding of the long-term impact hazard to Earth.

Acknowledgments

We thank MIT students Lindsey Malcom, Nancy Hsand April Deet Russell who were involved in various dprocessing or early analysis stages. R.P.B. acknowlesupport for this research by NASA Grant NAG5-12355 aNSF Grant AST-0205863with additional funding supporfrom The Planetary Society. We thank many colleagumost especially W. Bottke and A. Morbidelli, for their mahelpful discussions that helped shape the ideas and resupresented here. We are grateful to A. Morbidelli andClark for their supportive and helpful reviews. We thathe Bob Barr and the staff at the MDM Observatory whthis research originated. Binzel and Rivkin were VisitiAstronomers at Kitt Peak National Observatory, NatioOptical Astronomy Observatory, which is operated byAssociation of Universities for Research in Astronomy, I(AURA) under cooperative agreement with the National Sence Foundation. Binzel, Burbine, and Stuart had the plege to be Visiting Astronomers at the Infrared TelescFacility, which is operated by the University of Hawaii undCooperative Agreement no. NCC 5-538 with the NatioAeronautics and Space Administration, Office of Spaceence, Planetary Astronomy Program. Observations obtaat the Hale Telescope, Palomar Observatory are partcollaboration between the California Institute of Technogy, NASA/JPL, and Cornell University. The work at thePropulsion Laboratory, Caltech, was supported undertract from NASA. All newly reported spectral data presenhere are available athttp://smass.mit.edu/.

-

-

s

Appendix AObservation summary

Number and name Provisional designation Observing date Telesc

433 Eros 1898 DQ 2-Sep-95 MDM 1.3m433 Eros 1898 DQ 3-Sep-95 MDM 1.3m433 Eros 1898 DQ 27-Oct-95 MDM 1.3m433 Eros 1898 DQ 2-Dec-95 IRTF 3m433 Eros 1898 DQ 9-Dec-95 MDM 2.4m433 Eros 1898 DQ 29-Jan-96 MDM 2.4719 Albert 1911 MT 23-Oct-01 Palomar 5

1011 Laodamia 1924 PK 20-Jan-02 KPNO 4m1036 Ganymed 1924 TD 5-Nov-94 MDM 1.3m1036 Ganymed 1924 TD 24-Feb-97 MDM 2.41620 Geographos 1951 RA 7-Jan-94 MDM 2.41627 Ivar 1929 SH 9-May-95 MDM 2.4m1627 Ivar 1929 SH 10-Feb-97 IRTF 3m1627 Ivar 1929 SH 25-Feb-97 MDM 2.4m1862 Apollo 1932 HA 2-Dec-96 MDM 2.4m1862 Apollo 1932 HA 4-Jan-97 IRTF 3m1864 Daedalus 1971 FA 20-Feb-94 MDM 2.41865 Cerberus 1971 UA 14-Oct-98 MDM 2.41866 Sisyphus 1972 XA 6-Jan-94 MDM 2.41916 Boreas 1953 RA 23-Oct-01 Palomar 51917 Cuyo 1968 AA 1-Apr-94 MDM 2.4m1980 Tezcatlipoca 1950 LA 28-May-97 MDM 2.4m2062 Aten 1976 AA 9-Feb-95 MDM 2.4m2078 Nanking 1975 AD 2-Jan-97 IRTF 3m2078 Nanking 1975 AD 10-Feb-97 IRTF 3m2100 Ra-Shalom 1978 RA 4-Jan-97 IRTF 3m2100 Ra-Shalom 1978 RA 13-Sep-97 MDM 2.42100 Ra-Shalom 1978 RA 30-Sep-97 IRTF 3m2102 Tantalus 1975 YA 10-May-95 MDM 2.4m2102 Tantalus 1975 YA 12-May-95 MDM 2.4m2201 Oljato 1947 XC 10-Dec-95 MDM 2.4m2335 James 1974 UB 3-Mar-00 IRTF 3m2340 Hathor 1976 UA 23-Nov-97 MDM 2.4m2340 Hathor 1976 UA 24-Nov-97 MDM 2.4m2423 Ibarruri 1972 NC 3-Mar-00 IRTF 3m3102 Krok 1981 QA 6-May-00 IRTF 3m3102 Krok 1981 QA 6-Jul-00 Palomar 53103 Eger 1982 BB 31-Mar-94 MDM 2.4m3103 Eger 1982 BB 1-Apr-94 MDM 2.4m3103 Eger 1982 BB 10-Feb-97 IRTF 3m3103 Eger 1982 BB 25-Feb-97 MDM 2.4m3122 Florence 1981 ET3 22-Jan-97 MDM 2.43122 Florence 1981 ET3 10-Feb-97 IRTF 3m3199 Nefertiti 1982 RA 10-Feb-97 IRTF 3m3199 Nefertiti 1982 RA 24-Feb-97 MDM 2.4m3200 Phaethon 1983 TB 15-Nov-94 MDM 2.43288 Seleucus 1982 DV 24-Dec-01 Palomar3352 McAuliffe 1981 CW 22-Feb-94 MDM 2.4m3352 McAuliffe 1981 CW 27-Feb-99 IRTF 3m3552 Don Quixote 1983 SA 5-May-00 IRTF 3m3552 Don Quixote 1983 SA 23-Oct-01 Palomar 53671 Dionysus 1984 KD 7-Apr-97 MDM 2.4m3671 Dionysus 1984 KD 25-May-97 MDM 2.4m3674 Erbisbuhl 1963 RH 9-Aug-99 IRTF 3m3691 Bede 1982 FT 29-Jan-96 MDM 2.43753 Cruithne 1986 TO 13-Sep-97 MDM 2.43753 Cruithne 1986 TO 30-Sep-97 IRTF 3m3908 Nyx 1980 PA 9-Sep-96 MDM 2.4m3908 Nyx 1980 PA 12-Oct-96 MDM 2.4m3908 Nyx 1980 PA 4-Jan-97 IRTF 3m4034 1986 PA 6-Aug-97 Keck 10m4055 Magellan 1985 DO2 1-Mar-00 KPNO 4m




Appendix A (continued)

Number and name Provisional designation Observing date Telescope

4055 Magellan 1985 DO2 3-Mar-00 IRTF 3m4055 Magellan 1985 DO2 4-Mar-00 IRTF 3m4179 Toutatis 1989 AC 20-Jan-97 MDM 2.4m4183 Cuno 1959 LM 23-Nov-97 MDM 2.4m4183 Cuno 1959 LM 8-Jan-98 MDM 2.4m4197 1982 TA 11-Sep-96 MDM 2.4m4197 1982 TA 12-Oct-96 MDM 2.4m4341 Poseidon 1987 KF 25-May-97 MDM 2.4m4451 Grieve 1988 JJ 20-Jan-02 KPNO 4m4503 Cleobulus 1989 WM 25-Feb-99 KPNO 4m4688 1980 WF 16-Dec-00 Palomar 5m4947 Ninkasi 1988 TJ1 12-Oct-96 MDM 2.4m4954 Eric 1990 SQ 24-Feb-94 MDM 2.4m4954 Eric 1990 SQ 28-Mar-94 MDM 2.4m4954 Eric 1990 SQ 8-Feb-97 IRTF 3m4954 Eric 1990 SQ 2-Mar-00 IRTF 3m4957 Brucemurray 1990 XJ 1-Dec-96 MDM 2.4m5131 1990 BG 7-Feb-95 MDM 2.4m5131 1990 BG 27-Jan-99 IRTF 3m5143 Heracles 1991 VL 12-Oct-96 MDM 2.4m5275 Zdislava 1986 UU 10-Aug-99 IRTF 3m5587 1990 SB 17-May-01 KPNO 4m5604 1992 FE 1-Mar-00 KPNO 4m5626 1991 FE 19-Feb-94 MDM 2.4m5626 1991 FE 1-Apr-94 MDM 2.4m5641 McCleese 1990 DJ 2-Mar-00 IRTF 3m5646 1990 TR 28-Apr-96 MDM 2.4m5660 1974 MA 22-Aug-93 MDM 1.3m5751 Zao 1992 AC 17-Nov-94 MDM 2.4m5751 Zao 1992 AC 9-Feb-95 MDM 2.4m5828 1991 AM 22-Feb-02 KPNO 4m5836 1993 MF 21-Aug-93 MDM 1.3m5836 1993 MF 1-Oct-97 IRTF 3m6047 1991 TB1 17-Feb-98 IRTF 3m6047 1991 TB1 15-Oct-98 MDM 2.4m6047 1991 TB1 23-Mar-99 MDM 2.4m6053 1993 BW3 9-Dec-95 MDM 2.4m6455 1992 HE 8-Feb-95 MDM 2.4m6489 Golevka 1991 JX 9-May-95 MDM 2.4m6489 Golevka 1991 JX 22-May-99 IRTF 3m6569 1993 MO 10-May-95 MDM 2.4m6611 1993 VW 7-Jan-94 MDM 2.4m7336 Saunders 1989 RS1 11-Sep-96 MDM 2.4m7336 Saunders 1989 RS1 12-Oct-96 MDM 2.4m7341 1991 VK 12-Oct-96 MDM 2.4m7358 Oze 1995 YA3 29-Jan-96 MDM 2.4m7358 Oze 1995 YA3 15-Sep-98 IRTF 3m7480 Norwan 1994 PC 1-Dec-96 MDM 2.4m7482 1994 PC1 19-Jan-97 MDM 2.4m7822 1991 CS 24-Feb-97 MDM 2.4m7888 1993 UC 28-Mar-94 MDM 2.4m7889 1994 LX 2-May-98 IRTF 3m7889 1994 LX 16-May-01 KPNO 4m7977 1977 QQ5 8-Jan-98 MDM 2.4m8176 1991 WA 11-Dec-95 MDM 2.4m8566 1996 EN 28-Apr-96 MDM 2.4m9400 1994 TW1 14-Nov-94 MDM 2.4m9400 1994 TW1 7-Feb-95 MDM 2.4m

10115 1992 SK 27-Feb-99 IRTF 3m10165 1995 BL2 9-Feb-95 MDM 2.4m10563 Izhdubar 1993 WD 11-Dec-95 MDM 2.4m11311 Peleus 1993 XN2 5-Jan-94 MDM 2.4m11398 1998 YP11 24-Feb-99 KPNO 4m11398 1998 YP11 27-Feb-99 IRTF 3m


Number and name Provisional designation Observing date Telesco

11405 1999 CV3 25-Feb-99 KPNO 4m11405 1999 CV3 27-Feb-99 IRTF 3m11500 1989 UR 15-Oct-98 MDM 2.4m12538 1998 OH 7-May-00 IRTF 3m12711 1991 BB 7-Feb-96 MDM 2.4m12711 1991 BB 15-Jan-00 IRTF 3m12923 1999 GK4 23-May-99 IRTF 3m13651 1997 BR 8-Feb-97 IRTF 3m13651 1997 BR 25-Feb-97 MDM 2.4m14402 1991 DB 29-Feb-00 KPNO 4m14402 1991 DB 2-Mar-00 IRTF 3m15745 1991 PM5 6-Jul-00 Palomar 515817 Lucianotesi 1994 QC 6-Aug-97 Keck 10m16657 1993 UB 7-Jan-94 MDM 2.4m16960 1998 QS52 15-Oct-98 MDM 2.4m17274 2000 LC16 6-Jul-00 Palomar 517511 1992 QN 15-Dec-95 MDM 2.4m17511 1992 QN 12-Oct-96 MDM 2.4m18736 1998 NU 17-May-01 KPNO 4m19356 1997 GH3 10-Apr-97 MDM 2.4m20255 1998 FX2 28-Mar-98 MDM 2.4m20255 1998 FX2 30-Apr-98 IRTF 3m20425 1998 VD35 17-Dec-00 Palomar 520790 2000 SE45 16-Dec-00 Palomar 520826 2000 UV13 16-Dec-00 Palomar 522099 2000 EX106 5-May-00 IRTF 3m22771 1999 CU3 23-Oct-01 Palomar 523548 1994 EF2 30-Mar-94 MDM 2.4m23548 1994 EF2 1-Apr-94 MDM 2.4m24475 2000 VN2 16-Dec-00 Palomar 525330 1999 KV4 16-Dec-00 Palomar 531345 1998 PG 18-Sep-98 IRTF 3m31345 1998 PG 13-Oct-98 MDM 2.4m31346 1998 PB1 15-Sep-98 IRTF 3m32906 1994 RH 23-Feb-02 KPNO 4m35107 1991 VH 26-Feb-97 MDM 2.4m35670 1998 SU27 6-Mar-02 Palomar 536284 2000 DM8 22-Feb-02 KPNO 4m37336 2001 RM 27-Oct-01 KPNO 4m40310 1999 KU4 24-May-99 IRTF 3m48603 1995 BC2 9-Feb-95 MDM 2.4m53319 1999 JM8 11-Aug-99 IRTF 3m

1989 VA 15-Nov-94 MDM 2.4m1991 BN 28-Nov-02 Palomar 5m1992 BF 8-Jan-98 MDM 2.4m1993 TQ2 7-Jan-94 MDM 2.4m1994 AB1 20-Feb-94 MDM 2.4m1994 AW1 17-Dec-00 Palomar 5m1994 TF2 6-Aug-97 Keck 10m1995 WL8 14-Dec-95 MDM 2.4m1995 WL8 15-Dec-95 MDM 2.4m1996 BZ3 9-Feb-96 MDM 2.4m1996 FQ3 28-Apr-96 MDM 2.4m1996 FQ3 30-Apr-96 MDM 2.4m1997 AC11 18-Jan-97 MDM 2.4m1997 AQ18 18-Jan-97 MDM 2.4m1997 BQ 9-Feb-97 IRTF 3m1997 BQ 24-Feb-97 MDM 2.4m1997 CZ5 9-Feb-97 IRTF 3m1997 GL3 10-Apr-97 MDM 2.4m1997 RT 13-Sep-97 MDM 2.4m1997 RT 30-Sep-97 IRTF 3m1997 SE5 23-Nov-97 MDM 2.4m


pe

m

m

mmm

m

mm

m



Number and name Provisional designation Observing date Telescope

1997 TT25 23-Nov-97 MDM 2.4m1997 UH9 23-Nov-97 MDM 2.4m1997 US9 23-Nov-97 MDM 2.4m1998 FM5 28-Mar-98 MDM 2.4m1998 FM5 30-Apr-98 IRTF 3m1998 HT31 2-May-98 IRTF 3m1998 KU2 16-Sep-98 IRTF 3m1998 KU2 13-Oct-98 MDM 2.4m1998 MQ 9-Dec-98 MDM 2.4m1998 QR15 18-Sep-98 IRTF 3m1998 QR15 13-Oct-98 MDM 2.4m1998 SG2 15-Oct-98 MDM 2.4m1998 VO33 8-Dec-98 MDM 2.4m1998 VR 9-Dec-98 MDM 2.4m1998 WM 8-Dec-98 MDM 2.4m1998 WZ6 8-Dec-98 MDM 2.4m1998 XB 27-Jan-99 IRTF 3m1998 XM4 17-Dec-00 Palomar 5m1999 AQ10 27-Jan-99 IRTF 3m1999 CV8 24-Feb-99 KPNO 4m1999 CW8 24-Feb-99 KPNO 4m1999 DB2 24-Feb-99 KPNO 4m1999 DY2 24-Feb-99 KPNO 4m1999 EE5 23-Mar-99 MDM 2.4m1999 FB 23-Mar-99 MDM 2.4m1999 FK21 22-Feb-02 KPNO 4m1999 HF1 23-May-99 IRTF 3m1999 KW4 23-May-99 IRTF 3m1999 NC43 1-Mar-00 KPNO 4m1999 NC43 3-Mar-00 IRTF 3m2000 BG19 1-Mar-00 KPNO 4m2000 BG19 4-Mar-00 IRTF 3m2000 BG19 6-May-00 IRTF 3m2000 BJ19 17-Dec-00 Palomar 5m2000 CE59 1-Mar-00 KPNO 4m2000 CK33 21-Jan-02 KPNO 4m2000 CN33 1-Mar-00 KPNO 4m2000 CO101 1-Mar-00 KPNO 4m2000 CO101 2-Mar-00 IRTF 3m2000 DO1 1-Mar-00 KPNO 4m2000 DO1 3-Mar-00 IRTF 3m2000 DO8 3-Mar-00 IRTF 3m2000 EA107 5-Mar-02 KPNO 4m2000 EZ148 5-May-00 IRTF 3m2000 GD2 5-Mar-02 KPNO 4m2000 GJ147 5-May-00 IRTF 3m2000 GK137 6-Jul-00 Palomar 5m2000 GO82 7-May-00 IRTF 3m2000 GR146 5-May-00 IRTF 3m2000 GU127 6-May-00 IRTF 3m2000 GV127 5-May-00 IRTF 3m2000 JG5 7-May-00 IRTF 3m2000 JQ66 6-Jul-00 Palomar 5m2000 KL33 6-Jul-00 Palomar 5m2000 MU1 6-Jul-00 Palomar 5m2000 NM 6-Jul-00 Palomar 5m2000 OJ8 17-Dec-00 Palomar 5m2000 PG3 20-Jun-01 Magellan 6.5m2000 RW37 7-Mar-01 KPNO 4m2000 SY162 17-Dec-00 Palomar 5m2000 WC67 18-Jan-01 Palomar 5m2000 WF6 17-Dec-00 Palomar 5m2000 WJ10 17-Dec-00 Palomar 5m2000 WJ63 16-Dec-00 Palomar 5m


Number and name Provisional designation Observing date Telesc

2000 WK10 6-Mar-02 Palomar 5m2000 WL10 16-Dec-00 Palomar 52000 WL63 16-Dec-00 Palomar 52000 WM63 17-Dec-00 Palomar 52000 WO107 16-Dec-00 Palomar 52000 XL44 7-Mar-01 KPNO 4m2000 YA 17-Dec-00 Palomar 5m2000 YH66 18-Jan-01 Palomar 52000 YO29 18-Jan-01 Palomar 52001 DU8 7-Mar-01 KPNO 4m2001 EB 17-May-01 KPNO 4m2001 EC 7-Mar-01 KPNO 4m2001 FY 16-May-01 KPNO 4m2001 HA8 16-May-01 KPNO 4m2001 HK31 16-May-01 KPNO 4m2001 HW15 16-May-01 KPNO 4m2001 JM1 16-May-01 KPNO 4m2001 JV1 17-May-01 KPNO 4m2001 MF1 23-Dec-01 Palomar 52001 OE84 23-Oct-01 Palomar 52001 PD1 28-Oct-01 KPNO 4m2001 QA143 23-Oct-01 Palomar 52001 QQ142 6-Mar-02 KPNO 4m2001 SJ262 27-Oct-01 KPNO 4m2001 TC45 27-Oct-01 KPNO 4m2001 TX16 27-Oct-01 KPNO 4m2001 TY44 24-Dec-01 Palomar 52001 UA5 23-Feb-02 KPNO 4m2001 UC5 23-Oct-01 Palomar 52001 UU92 24-Dec-01 Palomar 52001 UY4 27-Oct-01 KPNO 4m2001 VG5 24-Dec-01 Palomar 52001 VS78 21-Jan-02 KPNO 4m2001 WA25 24-Dec-01 Palomar 52001 WG2 24-Dec-01 Palomar 52001 WH2 23-Dec-01 Palomar 52001 WL15 24-Dec-01 Palomar 52001 XN254 22-Feb-02 KPNO 4m2001 XR1 20-Jan-02 KPNO 4m2001 XS1 24-Dec-01 Palomar 52001 XS30 23-Dec-01 Palomar 52001 XU30 24-Dec-01 Palomar 52001 XY10 24-Dec-01 Palomar 5m2001 YE1 24-Dec-01 Palomar 52001 YK4 20-Jan-02 KPNO 4m2002 AA 21-Jan-02 KPNO 4m2002 AD9 21-Jan-02 KPNO 4m2002 AH29 21-Jan-02 KPNO 4m2002 AK14 20-Jan-02 KPNO 4m2002 AL14 6-Mar-02 Palomar 5m2002 AQ2 5-Mar-02 KPNO 4m2002 AU5 20-Jan-02 KPNO 4m2002 AV 20-Jan-02 KPNO 4m2002 BA1 6-Mar-02 KPNO 4m2002 BK25 22-Feb-02 KPNO 4m2002 BM26 22-Feb-02 KPNO 4m2002 BP26 23-Feb-02 KPNO 4m2002 CS11 5-Mar-02 KPNO 4m2002 CT46 22-Feb-02 KPNO 4m2002 DH2 5-Mar-02 KPNO 4m2002 DO3 6-Mar-02 Palomar 5m2002 DY3 5-Mar-02 KPNO 4m2002 EA 6-Mar-02 Palomar 5m2002 EC 5-Mar-02 KPNO 4m

ope

mmmm

mm

mm

m

m

mm

m

mmmm

mmm

m


Appendix BSpectra for near-Earth and Mars-crossing asteroids. These data are available in digital format athttp://smass.mit.edu/




Appendix B (continued)



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