Icarus 170 (2004) 259–294
www.elsevier.com/locate/icarus
Observed spectral properties of near-Earth objects: results for population
distribution, source regions, and space weathering processes
Richard P. Binzel a,∗ , Andrew S. Rivkin a , J. Scott Stuart a , Alan W. Harris b , Schelte J. Bus c ,
Thomas H. Burbine d
a Department of Earth, Atmospheric, and Planetary Sciences, Massachusetts Institute of Technology Cambridge, MA 02139, USA
b Space Science Institute, 4603 Orange Knoll, La Canada, CA 91011, USA
c Institute for Astronomy, 640 North A’ohoku Place, Hilo, HI 96720, USA
d Laboratory for Extraterrestrial Physics, NASA Goddard 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) objects observed
from 1994 through 2002 as a complement to the Small Main-Belt Asteroid Spectroscopic Survey (SMASS, http://smass.mit.edu/). Combined
with previously published SMASS results, we have an internally consistent data set of more than 400 of these objects for investigating trends
related to size, orbits, and dynamical history. These data also provide the basis for producing a bias-corrected estimate for the total NEO
population (Stuart and Binzel, 2004, Icarus 170, 295–311). We find 25 of the 26 Bus (1999, PhD thesis) taxonomic types are represented,
with nearly 90% of the objects falling within the broad S-, Q-, X-, and C-complexes. Rare A- and E-types are more common in the MC than
NEO population (about 5% compared to < 1%) and may be direct evidence of slow diffusion into MC orbits from the Flora and Hungaria
regions, respectively. A possible family of MC objects (C-types) may reside at the edge of the 5:2 jovian resonance. Distinct signatures are
revealed for the relative contributions of different taxonomic types to the NEO population through different source regions. E-types show
an 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 effectively
identical main-belt source region profiles, as would be expected if they have related origins. A lack of V-types among Mars-crossers suggests
entry into NEO space via rapid transport through the ν6 and 3:1 resonances from low eccentricity main-belt orbits, consistent with a Vesta
origin. D-types show the strongest signature from Jupiter family comets (JFC), with a strong JFC component also seen among the X-types.
A distinct taxonomic difference is found with respect to the jovian Tisserand parameter T , where C-, D-, and X-type (most likely low
albedo P-class) objects predominate for T 3. These objects, which may be extinct comets, comprise 4% of our observed sample, but their
low albedos makes this magnitude limited fraction under-representative of the true value. With our taxonomy statistics providing a strong
component to the diameter limited bias correction analysis of Stuart (2003, PhD thesis), we estimate 10–18% of the NEO population above
any given diameter may be extinct comets, taking into account asteroids scattered into T < 3 orbits and comets scattered into T > 3 orbits.
In terms of possible space weathering effects, we see a size-dependent transition from ordinary chondrite-like (Q-type) objects to S-type
asteroids over the size range of 0.1 to 5 km, where the transition is effectively complete at 5 km. A match between the average surface age
of 5 km asteroids and the rate of space weathering could constrain models for both processes. However, space weathering may proceed at
a very rapid rate compared with collisional timescales. In this case, the presence or absence of a regolith may be the determining factor for
whether or not an object appears “space weathered.” Thus 0.1 to 5 km appears to be a critical size range for understanding the processes,
timescales, 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
1. Introduction
* 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
Many pieces of the puzzle must be brought together in
order to have a clear picture of the near-Earth object (NEO)
population. Four of the pieces that can be described in-
260
R.P. Binzel et al. / Icarus 170 (2004) 259–294
clude: (i) the taxonomic distribution of the population as
measured by observational sampling, (ii) the determination
of albedos that can be associated with the taxonomic distribution, (iii) discovery statistics for the NEO population,
and (iv) the debiasing of the discovery statistics using the
taxonomic and albedo information. This paper presents the
first piece, detailing the observations and observed characteristics of the NEO and Mars-crossing (MC) population.
For the second piece, a complementary program of albedo
measurements was pursued at the Keck Observatory (Binzel,
P.I.) with results presented by Delbo et al. (2003). For the
third piece, the most extensive NEO discovery statistics
are provided by the LINEAR survey (Stokes et al., 2002;
Stuart, 2001). The work of Stuart (2003) brings the fourth
piece, which appears here as a companion paper by Stuart
and Binzel (2004).
The observations presented here were obtained as part
of the Small Main-Belt Asteroid Spectroscopic Survey
(SMASS) initiated at the Massachusetts Institute of Technology in the early 1990’s. SMASS was undertaken to extend
our knowledge of the compositional properties of main-belt
asteroids to smaller and smaller sizes by taking advantage
of state of the art charged–coupled device (CCD) detectors.
The first stage of this survey, which we herein refer to as
SMASSI, is reported by Xu et al. (1995). The second stage,
SMASSII, is reported by Bus and Binzel (2002a). NearEarth objects, which we define as belonging to the Aten,
Apollo, and Amor groups generally having perihelion distances less than 1.3 AU from the Sun (Shoemaker et al.,
1979), and Mars-crossing objects have been routinely observed as targets of opportunity within the SMASS program
since its inception. As the SMASSII program achieved its
goals (more than 1300 main-belt asteroids measured and an
extended system of asteroid taxonomy; Bus, 1999; Bus and
Binzel, 2002a, 2002b), increasing focus has been directed
toward the measurement of NEOs. This focus has been enabled by the increasing discovery rate and availability of
NEO targets for measurement (Stokes et al., 2002).
The fundamental scientific rationale for sampling the
physical properties of near-Earth objects is reviewed by
Binzel et al. (2002). NEOs are “immigrants” to the inner Solar System with lifetimes of order 106 –107 years (Morbidelli
et al., 2002a). Because these lifetimes are extremely short
compared to the age of the Solar System, the current population must be re-supplied. We seek to understand the origin
of the NEO population from both asteroidal and possible
cometary sources. What’s more, meteorites are (by definition of their orbital intersection) near-Earth objects prior
to their arrival—thus correlations between NEOs and meteorites represent the most direct link for achieving an understanding of asteroid–meteorite relationships. In addition,
NEOs are the smallest telescopically observable Solar System bodies. By using NEOs to leverage over the greatest
possible size range, we seek to investigate possible compositional trends as a function of size that may be indicative of
processes such as space weathering that modify the reflective properties of asteroid surfaces.
By virtue of their proximity, NEOs are among the most
accessible spacecraft destinations in our Solar System. Sample return missions (e.g., Sears et al., 2000) have the potential to address these science objectives (and more) using
the full capabilities of ground based analysis facilities. Reconnaissance of NEOs to select the most scientifically rewarding candidates for sample return missions is a further
motivation for the SMASS NEO observations (Binzel et al.,
2004b). Moving from the scientific to the pragmatic, assessing the NEO impact hazard (Morrison et al., 2002) requires
knowledge of the compositional distribution of the population. From the compositional distribution, albedo models
can be applied to convert the discovery H magnitude distribution into an actual size distribution (Delbo et al., 2003;
Stuart, 2003; Stuart and Binzel, 2004). Similarly, density assumptions can be applied to taxonomic categories to yield
mass estimates for impact energy distributions. Binzel et al.
(2003) argue that at this stage of our knowledge, the measurement goals for scientifically understanding the NEOs are
the same as the measurement goals for improving our assessment of their impact hazard.
We organize this paper as follows. In Section 2 we present
a description and compilation of the SMASS NEO observations with the newly reported spectra being displayed in the
Appendix. Section 3 presents a taxonomic analysis of the
SMASS NEO data as well as other published NEO results.
Throughout this paper we utilize the Bus (1999) taxonomy,
with the addition of albedo data (where available) to distinguish E-, M-, and P-types. These taxonomic results provide
basic input to the debiased population model derived in our
companion paper (Stuart and Binzel, 2004). Section 4 looks
at the taxonomic and dynamic traceability of NEOs to their
source regions while Section 5 reveals new trends that may
be related to space weathering. Section 6 presents a discussion of these trends and implications for the cometary
contribution to the NEO population. Concluding remarks are
given in Section 7.
2. SMASS observations
Table 1 provides a summary compilation of SMASS results for 401 near-Earth and Mars-crossing objects. Through
the application of similar observing and reduction procedures on a consistent set of telescopes, this table provides the
largest available uniform data set for this population. Within
the table, MC, ATE, APO, AMO denotes the orbital class
as Mars-crossing, Aten, Apollo, or Amor. The taxonomic
type, Slope, and principal component (PC) scores follow
the system of Bus (1999), as discussed in Section 3. The
absolute (H) magnitudes are from the Minor Planet Center
database. The final column of Table 1 includes references
for the objects whose SMASS spectra have been previously
published. New SMASS spectra for 252 of these entries are
Spectral properties of near-Earth objects
261
Table 1
SMASS observational results for near-Earth and Mars-crossing objects
Number and name
132
391
433
512
699
719
985
1011
1036
1065
1131
1134
1139
1198
1204
1293
1316
1374
1565
1593
1620
1627
1640
1660
1685
1862
1863
1864
1865
1866
1916
1917
1943
1951
1980
1981
2035
2062
2063
2064
2074
2078
2099
2100
2102
2201
2204
2253
2335
2340
2423
2629
2744
3040
3102
3103
3122
3198
3199
3200
3216
3255
Aethra
Ingeborg
Eros
Taurinensis
Hela
Albert
Rosina
Laodamia
Ganymed
Amundsenia
Porzia
Kepler
Atami
Atlantis
Renzia
Sonja
Kasan
Isora
Lemaitre
Fagnes
Geographos
Ivar
Nemo
Wood
Toro
Apollo
Antinous
Daedalus
Cerberus
Sisyphus
Boreas
Cuyo
Anteros
Lick
Tezcatlipoca
Midas
Stearns
Aten
Bacchus
Thomsen
Shoemaker
Nanking
Opik
Ra-Shalom
Tantalus
Oljato
Lyyli
Espinette
James
Hathor
Ibarruri
Rudra
Birgitta
Kozai
Krok
Eger
Florence
Wallonia
Nefertiti
Phaethon
Harrington
Tholen
Provisional designation
Orbit
Type
1953 LF
1934 AJ
1898 DQ
1903 LV
1957 WX1
1911 MT
1922 MO
1924 PK
1924 TD
1955 SM1
1929 RO
1951 SA
1929 XE
1931 RA
1931 TE
1933 SO
1978 WK14
1935 UA
1948 WA
1951 LB
1951 RA
1929 SH
1951 QA
1953 GA
1948 OA
1932 HA
1948 EA
1971 FA
1971 UA
1972 XA
1953 RA
1968 AA
1973 EC
1949 OA
1950 LA
1973 EA
1973 SC
1976 AA
1977 HB
1942 RQ
1974 UA
1975 AD
1977 VB
1978 RA
1975 YA
1947 XC
1943 EQ
1932 PB
1974 UB
1976 UA
1972 NC
1980 RB1
1975 RB
1979 BA
1981 QA
1982 BB
1981 ET3
1981 YH1
1982 RA
1983 TB
1980 RB
1980 RA
MC
MC
AMO
MC
MC
AMO
MC
MC
AMO
MC
MC
MC
MC
MC
MC
MC
MC
MC
MC
MC
APO
AMO
MC
MC
APO
APO
APO
APO
APO
APO
AMO
AMO
AMO
MC
AMO
APO
MC
ATE
APO
MC
MC
MC
MC
ATE
APO
APO
MC
MC
MC
ATE
MC
MC
MC
MC
AMO
APO
AMO
MC
AMO
APO
MC
MC
Xe
S
S
S
Sq
S
S
S
S
S
S
S
S
L
S
Sq
Sr
Sq
Sq
S
S
S
S
S
S
Q
Sq
Sr
S
S
S
Sl
L
A
Sl
V
Xe
Sr
Sq
S
Sa
Sq
Ch
C
Q
Sq
X
Sl
Sa
Sq
A
B
S
S
S
Xe
S
S
Sq
B
S
S
Slope
0.1605
0.5312
0.6271
0.4488
0.2152
0.4553
0.5062
0.5650
0.4298
0.5563
0.5324
0.5075
0.4885
0.7604
0.4704
0.2298
0.2799
0.2313
0.1831
0.5191
0.3892
0.5616
0.4981
0.4426
0.3031
0.0658
0.1776
0.2560
0.3085
0.6488
0.5044
0.7233
0.6400
1.1774
0.6492
−0.3096
0.4441
0.3008
0.2484
0.4810
0.8070
0.2431
−0.0292
0.0765
−0.0186
−0.0055
0.3596
0.7030
0.6903
0.0888
0.8595
−0.2512
0.3820
0.5770
0.3911
0.5988
0.3136
0.4586
0.1479
−0.1941
0.4936
0.3313
PC2
PC3
H Mag
0.2078
−0.1769
−0.2581
−0.2840
−0.0519
−0.1187
−0.2221
−0.2118
−0.2054
−0.2377
−0.2669
−0.1024
−0.1872
0.0984
−0.1047
−0.2789
−0.3933
−0.1748
−0.2210
−0.3458
−0.1177
−0.3029
−0.0457
−0.1498
−0.2011
−0.3680
−0.1793
−0.3072
−0.1508
−0.1553
−0.1420
−0.2161
−0.0160
−0.5920
−0.2832
−0.7383
0.1200
−0.3810
−0.2110
−0.1581
−0.1836
−0.1220
0.3217
0.1810
−0.3202
−0.1095
0.3154
−0.0634
−0.3627
−0.0527
−0.3971
0.3722
−0.2635
−0.3017
−0.3109
0.0946
−0.2703
−0.0100
−0.0205
0.3243
−0.1840
−0.2011
0.0063
0.0150
−0.0048
−0.0001
0.0129
−0.0205
−0.0030
−0.0421
0.0474
0.0291
−0.0055
0.0042
−0.0021
0.0827
−0.0037
0.0069
−0.0305
−0.0445
−0.0269
−0.0317
0.0190
−0.0515
0.0028
0.0310
−0.0146
−0.0251
0.0521
−0.0780
−0.0047
−0.0302
−0.0168
0.0105
−0.0690
−0.0676
−0.0045
0.0470
0.0128
−0.0785
−0.0685
−0.0701
0.0649
0.0117
−0.0673
0.0803
−0.0192
0.0441
0.0434
0.0603
−0.0020
0.0287
−0.0363
0.1447
−0.0435
−0.0152
−0.0513
−0.0371
0.0243
0.0954
−0.0353
0.0460
0.0571
0.0300
9.4
10.1
11.2
10.7
11.7
15.8
12.7
12.7
9.5
13.2
13.0
14.3
12.5
14.6
12.2
12.0
13.3
13.5
12.3
13.2
16.5
13.2
13.1
11.9
14.0
16.3
15.8
15.0
17.0
13.0
15.0
13.9
16.0
14.7
14.0
15.2
12.6
17.1
17.1
13.1
14.0
12.1
15.2
16.1
16.2
16.9
12.7
12.9
13.8
19.2
13.2
14.5
14.8
14.5
15.6
15.4
14.2
12.3
15.1
14.3
14.0
13.6
T
Reference
3.18
2
3.41
2
4.58
∗
3.62
2
3.24
2
3.14
∗
3.54
2
3.44
∗
3.03
∗
3.48
2
3.59
2
3.17
2
3.82
2
3.55
1
3.56
2
3.59
2
3.34
2
3.57
2
3.36
2
3.57
2
5.07
∗
3.88
∗
3.51
2
3.38
2
4.71
6
4.41
∗
3.30
3
4.33
∗
5.59
∗
3.51
∗
3.44
∗
3.43
∗
4.64
6
4.54
2
3.99
∗
3.61
3
3.82
2
6.18
∗
5.67
5
3.60
2
3.90
1
3.37
∗
3.36
2
6.94
∗
4.45
∗
3.30
∗
3.22
1
3.55
1
3.41
∗
6.88
∗
3.62
∗
4.02
2
3.51
2
3.63
2
3.55
∗
4.61
∗
3.92
∗
3.58
2
4.19
∗
4.51
∗
3.46
2
3.36
2
(continued on next page)
262
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Table 1 (continued)
Number and name
3287
3288
3352
3401
3416
3443
3552
3581
3635
3671
3674
3691
3737
3753
3800
3833
3858
3873
3908
3920
4034
4055
4142
4179
4183
4197
4205
4222
4276
4341
4435
4451
4503
4558
4660
4688
4910
4947
4954
4957
4995
5038
5131
5143
5230
5253
5275
5349
5392
5407
5510
5585
5587
5604
5626
5641
5646
5649
5660
5732
5751
5817
5828
Olmstead
Seleucus
McAuliffe
Vanphilos
Dorrit
Leetsungdao
Don Quixote
Alvarez
Dionysus
Erbisbuhl
Bede
Beckman
Cruithne
Karayusuf
Calingasta
Dorchester
Roddy
Nyx
Aubignan
Magellan
Dersu-Uzala
Toutatis
Cuno
David Hughes
Nancita
Clifford
Poseidon
Holt
Grieve
Cleobulus
Janesick
Nereus
Kawasato
Ninkasi
Eric
Brucemurray
Overbeek
Heracles
Asahina
Zdislava
Paulharris
Parker
Parks
McCleese
Donnashirley
Zao
Robertfrazer
Provisional designation
Orbit
Type
Slope
PC2
PC3
H Mag
1981 DK1
1982 DV
1981 CW
1981 PA
1931 VP
1979 SB1
1983 SA
1985 HC
1981 WO1
1984 KD
1963 RH
1982 FT
1983 PA
1986 TO
1984 AB
1971 SC
1986 TG
1984 WB
1980 PA
1948 WF
1986 PA
1985 DO2
1981 KE
1989 AC
1959 LM
1982 TA
1985 YP
1988 EK1
1981 XA
1987 KF
1983 AG2
1988 JJ
1989 WM
1988 NF
1982 DB
1980 WF
1953 PR
1988 TJ1
1990 SQ
1990 XJ
1984 QR
1948 KF
1990 BG
1991 VL
1988 EF
1985 XB
1986 UU
1988 RA
1986 AK
1992 AX
1988 RF7
1990 MJ
1990 SB
1992 FE
1991 FE
1990 DJ
1990 TR
1990 WZ2
1974 MA
1988 WC
1992 AC
1989 RZ
1991 AM
MC
AMO
AMO
MC
MC
MC
AMO
MC
MC
AMO
MC
AMO
MC
ATE
MC
MC
MC
MC
AMO
MC
APO
AMO
MC
APO
APO
APO
MC
MC
MC
APO
MC
MC
AMO
MC
APO
AMO
MC
AMO
AMO
AMO
MC
MC
APO
APO
MC
MC
MC
MC
MC
MC
MC
MC
AMO
ATE
AMO
MC
AMO
MC
APO
MC
AMO
MC
APO
L
K
A
S
Sa
T
D
B
S
Cb
Sk
Xc
S
Q
S
Cb
Sa
S
V
Sa
O
V
A
Sk
Sq
Sq
Xe
S
Cb
O
S
S
Sq
S
Xe
V
S
Sq
S
S
S
S
S
O
S
S
Sa
C
Sl
Sk
S
Ch
Sq
V
S
A
U
S
Q
S
X
S
Q
0.6200
0.4974
0.8718
0.4994
0.8073
0.5491
1.1147
−0.3237
0.4567
−0.0630
0.2658
0.2058
0.5326
0.0130
0.5372
−0.0195
0.6807
0.4203
−0.5268
0.6048
−0.1440
−0.2210
1.0608
0.2733
0.0019
0.1727
0.2392
0.3205
0.0977
−0.1709
0.4085
0.5270
0.1184
0.3118
0.2958
0.0852
0.5497
0.1981
0.5589
0.6394
0.3525
0.4931
0.5089
−0.1949
0.4305
0.6193
0.6226
−0.0416
0.6551
0.2271
0.6165
−0.0644
0.2478
−0.1308
0.5842
0.8533
−0.1015
0.3343
0.1450
0.4511
0.2418
0.5106
−0.0250
0.0130
−0.0240
−0.3484
−0.1016
−0.3391
0.1168
0.4581
0.3113
−0.2039
0.2376
−0.1008
0.0798
−0.1470
−0.2254
−0.3416
0.2452
−0.5277
−0.0952
−0.9300
−0.3837
−0.3152
−1.1119
−0.2850
−0.0239
−0.1273
−0.0500
0.1707
−0.2770
0.3230
−0.1810
−0.2522
−0.1972
−0.1224
−0.2297
0.2634
−0.4530
−0.2466
−0.2952
−0.2375
−0.0886
−0.1473
−0.1058
−0.1387
−0.1147
−0.2170
−0.2635
−0.4940
0.2833
−0.2174
−0.0085
−0.2987
0.3847
−0.0442
−1.2293
−0.2294
−0.3894
−0.6302
−0.2600
−0.3748
−0.1602
0.1032
−0.3254
−0.1546
−0.0081
−0.0790
−0.0783
0.0296
−0.0620
−0.0111
0.0703
0.0375
0.0788
0.0444
0.0235
−0.0490
0.0071
0.0731
−0.0440
0.0406
−0.1248
0.0663
0.0429
−0.0658
−0.0801
0.0336
−0.0009
0.0074
−0.0165
−0.0245
0.0168
−0.0435
0.1000
−0.0576
0.0096
0.0169
0.0049
0.0299
0.0409
−0.0582
−0.0053
−0.0278
0.0544
−0.0466
0.0105
−0.0213
0.0031
−0.0371
−0.0127
−0.0329
−0.0639
0.0030
−0.0876
0.0099
−0.0099
0.0111
0.3496
0.0065
0.0263
−0.0773
−0.0859
0.0041
−0.1204
0.0177
0.0272
0.0172
−0.0113
15.0
15.3
15.8
12.6
13.7
13.3
13.0
12.1
14.8
16.7
12.1
14.9
13.0
15.1
15.4
15.0
13.7
12.0
17.4
13.2
18.1
14.8
13.6
15.3
14.4
14.9
14.7
12.4
14.3
15.5
13.2
12.2
15.6
12.2
18.3
19.0
14.2
18.7
12.6
15.1
13.0
14.1
14.1
14.0
13.4
13.7
13.6
12.7
12.7
13.7
13.9
13.7
13.6
16.4
14.7
12.7
16.1
13.2
15.7
14.1
14.9
12.6
16.3
T
Reference
3.46
2
3.66
∗
3.88
∗
3.37
2
3.81
2
3.43
2
2.31
∗
3.04
2
4.00
2
3.43
∗
3.37
∗
3.98
∗
3.33
2
5.92
∗
4.36
2
3.54
2
3.62
2
3.85
2
3.78
∗
3.56
2
5.70
∗
3.88
∗
3.79
6
3.15
∗
3.57
∗
3.09
∗
4.10
2
3.48
2
3.72
2
3.69
∗
3.41
2
3.15
∗
3.15
∗
3.49
2
4.49
6
3.44
∗
3.42
2
4.77
∗
3.66
∗
4.20
∗
3.41
2
3.51
2
4.21
∗
3.58
∗
3.35
2
3.69
2
3.62
∗
3.00
2
3.38
2
3.95
2
3.63
2
3.08
2
3.25
∗
6.38
∗
3.52
∗
3.94
∗
3.57
∗
3.44
2
3.51
∗
3.45
2
3.58
∗
3.35
2
3.77
∗
(continued on next page)
Spectral properties of near-Earth objects
263
Table 1 (continued)
Number and name
Provisional designation
Orbit
Type
Slope
PC2
PC3
H Mag
5836
6047
6053
6249
6386
6455
6489
6500
6569
6585
6611
6847
7304
7336
7341
7358
7474
7480
7482
7604
7753
7822
7888
7889
7977
8176
8201
8566
9400
9969
10115
10165
10302
10563
11066
11311
11398
11405
11500
12538
12711
12923
13651
14402
15745
15817
16064
16657
16834
16960
17274
17511
18514
18736
18882
19356
20043
20255
20425
20790
20826
22099
22449
1993 MF
1991 TB1
1993 BW3
1991 JF1
1989 NK1
1992 HE
1991 JX
1993 ET
1993 MO
1984 SR
1993 VW
1977 RL
1994 AE2
1989 RS1
1991 VK
1995 YA3
1992 TC
1994 PC
1994 PC1
1995 QY2
1988 XB
1991 CS
1993 UC
1994 LX
1977 QQ5
1991 WA
1994 AH2
1996 EN
1994 TW1
1992 KD
1992 SK
1995 BL2
1989 ML
1993 WD
1992 CC1
1993 XN2
1998 YP11
1999 CV3
1989 UR
1998 OH
1991 BB
1999 GK4
1997 BR
1991 DB
1991 PM5
1994 QC
1999 RH27
1993 UB
1997 WU22
1998 QS52
2000 LC16
1992 QN
1996 TE11
1998 NU
1999 YN4
1997 GH3
1993 EM
1998 FX2
1998 VD35
2000 SE45
2000 UV13
2000 EX106
1996 VC
AMO
APO
AMO
MC
MC
APO
APO
MC
AMO
MC
APO
MC
MC
AMO
APO
AMO
AMO
AMO
APO
MC
APO
APO
APO
APO
AMO
APO
APO
APO
AMO
MC
APO
APO
AMO
APO
APO
APO
AMO
APO
APO
APO
APO
APO
APO
AMO
AMO
AMO
AMO
AMO
APO
APO
AMO
APO
MC
AMO
AMO
AMO
MC
AMO
APO
AMO
APO
APO
MC
S
S
Sq
Xe
S
S
Q
B
Sr
Sk
V
Sk
Ld
Sq
Sq
Sq
X
S
S
C
B
S
U
V
S
Q
O
U
Sr
Q
S:
L
X
Q
K
Sq
Sr
Sq
S
S:
Sr
S:
S
C
S
Xc
C
Sr
S
Sq
Xk
X
Xc
Sk
S
S
U
Sq
Sq
S
Sq
S:
S
0.4099
0.5319
0.2602
0.4044
0.5220
0.4823
0.2310
−0.2539
0.4724
0.2748
−0.1473
0.2590
0.8737
−0.0671
0.0632
0.1666
0.3486
0.3011
0.4594
−0.0702
−0.2817
0.4548
1.0453
−0.4210
0.3087
0.0078
−0.1989
−0.2108
0.4645
0.0863
−0.1464
−0.1278
−0.2218
0.1247
−0.1874
−0.3305
−0.3507
0.3066
−0.4910
−0.1441
−0.7528
−0.0811
0.0157
−0.1819
−0.2138
−0.1383
0.3191
−0.2670
−0.2495
0.1381
0.3549
−0.1713
−0.7483
−1.2968
−0.1537
−0.2064
−0.2479
−1.6055
−0.4771
−0.2935
0.0222
−0.0136
0.0279
0.0499
0.0601
0.1027
−0.0482
0.0620
−0.1319
−0.0192
0.1062
−0.0009
0.1074
0.0018
−0.0211
0.0000
0.0702
0.0060
−0.0101
0.0179
0.1009
−0.0176
−0.2914
0.1178
0.0584
−0.0218
0.0195
0.1226
−0.0901
−0.0896
0.6313
0.1365
0.0549
0.4980
0.2460
0.2792
0.2284
0.3906
−0.0295
0.2478
−0.1667
0.0917
−0.1192
−0.3033
−0.2422
−0.1247
−0.0181
0.0443
−0.0073
0.0992
0.0050
−0.0382
−0.0182
−0.0103
0.3301
−0.3537
−0.0808
0.3346
0.0077
0.5345
0.1489
0.1197
0.2720
0.5482
−0.0205
0.3726
0.1482
0.1238
0.2540
0.3034
0.3065
−0.1013
0.2655
0.1093
0.3174
0.2232
−0.2565
0.2024
−0.1150
0.0752
0.2540
−0.2628
−0.0848
−0.0503
0.2333
0.2297
0.1468
−0.0260
−0.1712
−0.2516
0.0583
−0.1376
−0.3739
−0.0482
−0.1018
−0.1260
0.0214
0.0078
−0.1092
0.0044
0.0622
−0.0189
−0.0139
0.0021
0.0308
−0.0007
0.0737
0.0164
−0.0467
0.0576
0.0101
−0.0595
−0.0092
0.0049
0.3817
−0.1729
−0.0513
15.0
17.0
15.2
12.4
12.7
13.8
19.1
12.5
16.2
14.3
16.5
13.7
13.3
18.7
16.7
14.4
18.0
17.5
16.8
13.7
18.6
17.4
15.3
15.3
15.4
17.1
16.3
16.5
14.8
15.8
17.0
17.1
19.5
17.3
15.0
16.5
16.3
15.0
18.4
16.1
16.0
16.1
17.6
18.4
17.8
18.6
16.9
16.9
15.7
14.3
16.7
17.1
15.8
16.1
16.3
17.1
15.4
18.2
20.4
16.6
13.5
18.0
13.8
Jennifer
Golevka
Kodaira
O’Keefe
Kunz-Hallstei
Namiki
Saunders
Oze
Norwan
Braille
Izhdubar
Sigurd
Peleus
Lucianotesi
T
Reference
3.28
∗
4.48
∗
3.44
∗
3.78
2
3.54
2
3.18
∗
3.18
∗
3.04
2
4.21
∗
3.36
2
4.05
∗
3.40
2
3.25
2
3.41
∗
3.84
∗
3.49
∗
4.36
1
4.34
∗
4.66
∗
2.86
2
4.47
5
5.36
∗
3.05
∗
4.86
∗
3.38
∗
3.95
∗
3.02
3
4.22
∗
2.94
∗
3.28
3
5.06
∗
4.98
∗
5.06
3
5.54
∗
4.50
1
3.43
∗
4.05
∗
4.46
∗
5.65
∗
4.28
∗
5.10
∗
3.72
∗
4.81
∗
4.06
∗
4.10
∗
4.90
∗
3.02
3
3.35
∗
4.46
3
3.00
∗
3.10
∗
5.25
∗
3.15
2
3.38
∗
3.97
3
3.22
∗
3.81
3
3.52
∗
4.28
∗
3.09
∗
3.04
∗
5.58
∗
3.39
2
(continued on next page)
264
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Table 1 (continued)
Number and name
Provisional designation
Orbit
Type
Slope
PC2
PC3
H Mag
22771
23548
24475
25143 Itokawa
25330
26209
31345
31346
32906
35107
35396
35432
35670
36017
36183
36284
37336
40267
40310
47035
48603
53319
1999 CU3
1994 EF2
2000 VN2
1998 SF36
1999 KV4
1997 RD1
1998 PG
1998 PB1
1994 RH
1991 VH
1997 XF11
1998 BG9
1998 SU27
1999 ND43
1999 TX16
2000 DM8
2001 RM
1999 GJ4
1999 KU4
1998 WS
1995 BC2
1999 JM8
1989 UQ
1989 VA
1991 BN
1991 XB
1992 BF
1992 NA
1992 UB
1993 TQ2
1994 AB1
1994 AW1
1994 TF2
1995 BM2
1995 WL8
1995 WQ5
1996 BZ3
1996 FG3
1996 FQ3
1996 GT
1996 UK
1997 AC11
1997 AQ18
1997 BQ
1997 CZ5
1997 GL3
1997 RT
1997 SE5
1997 TT25
1997 UH9
1997 US9
1998 BB10
1998 BM10
1998 BT13
1998 FM5
1998 HT31
1998 KU2
1998 MQ
1998 MW5
1998 QR15
1998 SG2
1998 ST49
1998 UT18
APO
AMO
AMO
APO
APO
MC
AMO
AMO
AMO
APO
APO
AMO
APO
AMO
AMO
APO
AMO
APO
MC
MC
AMO
APO
ATE
ATE
APO
AMO
ATE
AMO
AMO
AMO
AMO
AMO
ATE
MC
AMO
MC
AMO
APO
AMO
AMO
MC
ATE
APO
APO
MC
APO
AMO
AMO
AMO
ATE
APO
APO
MC
APO
AMO
APO
AMO
AMO
APO
AMO
AMO
APO
APO
Sl
Q
Sa
S(IV)
B
Sq
Sq
Sq:
S
Sk
Xk
S
Sq
Sl
Ld
Sq
S
Sq
S:
Sr
X
X:
B
Sq
Q
K
Xc
C
X
Sa
Sq
L
Sr
Sq
Sq
Ch
X
C
Sq
Xk
Sq
Xc
C
S
S
V
O
T
Sq
Sq
Q
Sq
Sq
Sq
S
C:
Cb
S
Sq
Sq
Sq
Q
C
0.6226
0.1469
0.7374
0.4644
−0.0545
0.2589
0.1485
−0.0570
−0.3157
−0.2446
−0.3073
0.2652
−0.1600
−0.2232
−0.0364
0.0033
−0.0034
−0.2203
0.0129
−0.0158
−0.0576
0.2340
0.3109
0.2616
0.3359
0.0195
0.6535
0.9304
0.2411
0.2976
0.0082
−0.0092
−0.1960
0.0356
−0.1532
−0.0845
−0.1569
−0.0633
−0.1363
−0.0851
−0.2570
0.0012
−0.1196
−0.0273
0.0295
−0.0129
0.0989
0.0375
−0.0095
−0.0117
0.0999
0.2595
0.1629
−0.2472
0.1255
0.0605
0.0213
−0.0991
−0.0155
−0.0125
0.3583
0.1581
−0.0687
0.3485
0.6672
0.2612
0.6599
0.3730
0.2142
0.0950
−0.1551
0.1665
−0.0786
0.2013
0.4066
0.1050
0.0658
−0.0333
0.6060
0.3600
−0.3161
−0.1940
0.6256
0.2395
0.1526
−0.0392
0.1320
0.2596
0.1269
0.4468
0.2898
−0.2709
−0.1486
0.1423
0.0422
0.3189
0.2866
−0.2869
−0.1357
0.0345
−0.3147
−0.2681
−0.0616
0.3241
0.1918
0.2497
−0.1803
0.1959
−0.1322
0.1867
0.2344
−0.2500
−0.2067
−0.4887
−0.2061
0.2200
−0.1453
−0.1328
−0.1906
−0.1429
−0.1225
−0.1703
−0.2091
0.0625
−0.0440
−0.0345
0.0745
−0.0434
0.0517
0.0941
0.0048
0.0421
−0.0630
−0.1242
0.0153
0.0655
0.0139
0.0363
−0.0122
0.0854
−0.0164
0.0333
0.0106
0.0185
−0.0587
0.0000
0.1259
0.0186
0.0866
0.0063
−0.0146
−0.0305
0.0808
0.0496
0.0592
0.0017
−0.0044
0.3283
0.0694
0.0048
−0.0695
−0.0904
−0.0964
0.2611
−0.3191
−0.3656
−0.1452
−0.1749
−0.2829
0.1721
0.0877
0.0127
−0.0544
−0.0547
−0.0535
0.1064
−0.0078
17.0
17.6
16.5
19.2
16.8
16.0
17.6
17.1
16.0
16.9
16.9
19.5
19.6
19.2
16.2
14.9
15.2
15.0
16.8
12.5
17.3
15.2
19.0
17.9
19.3
18.1
19.5
16.5
16.0
20.0
16.3
17.7
19.3
15.2
18.1
16.8
18.2
17.8
21.0
18.5
16.4
21.0
18.2
18.0
13.6
20.0
20.0
14.8
19.3
18.8
17.3
20.4
16.5
26.5
16.0
20.8
16.6
16.6
19.2
18.0
19.7
17.7
20.4
T
Reference
4.22
∗
3.31
∗
3.70
∗
4.90
4
4.35
∗
3.23
2
3.72
∗
3.68
∗
3.43
∗
5.47
∗
4.53
5
3.21
3
3.47
∗
4.44
3
4.16
3
4.11
∗
3.23
∗
4.38
3
3.55
∗
3.12
2
3.80
∗
2.99
∗
6.49
5
7.66
∗
4.57
∗
2.93
1
6.52
∗
3.28
1
2.90
1
3.73
∗
3.02
∗
5.55
∗
6.00
∗
3.43
2
3.32
∗
3.36
2
3.18
∗
5.78
3
3.66
∗
4.20
5
3.19
2
6.36
∗
5.33
∗
3.98
∗
3.37
∗
3.10
∗
3.43
∗
2.66
∗
3.60
∗
6.90
∗
5.75
∗
4.97
3
3.31
3
3.19
3
3.37
∗
3.05
∗
3.40
∗
3.89
∗
5.90
3
3.07
∗
3.48
∗
3.23
3
3.01
3
(continued on next page)
Spectral properties of near-Earth objects
265
Table 1 (continued)
Number and name
Provisional designation
Orbit
Type
Slope
PC2
PC3
H Mag
1998 VO33
1998 VR
1998 WM
1998 WP5
1998 WZ6
1998 XB
1998 XM4
1999 AQ10
1999 CF9
1999 CV8
1999 CW8
1999 DB2
1999 DJ4
1999 DY2
1999 EE5
1999 FA
1999 FB
1999 FK21
1999 FN19
1999 HF1
1999 JD6
1999 JE1
1999 JO8
1999 JU3
1999 JV3
1999 JV6
1999 KW4
1999 NC43
1999 RB32
1999 SE10
1999 SK10
1999 VM40
1999 VN6
1999 VQ5
1999 WK13
1999 XO35
1999 YB
1999 YD
1999 YF3
1999 YG3
1999 YK5
2000 AC6
2000 AE205
2000 AH205
2000 AX93
2000 BG19
2000 BJ19
2000 BM19
2000 CE59
2000 CK33
2000 CN33
2000 CO101
2000 DO1
2000 DO8
2000 EA107
2000 EZ148
2000 GD2
2000 GJ147
2000 GK137
2000 GO82
2000 GR146
2000 GU127
2000 GV127
APO
ATE
APO
AMO
APO
ATE
APO
ATE
APO
APO
APO
AMO
APO
AMO
AMO
APO
APO
ATE
AMO
ATE
ATE
APO
AMO
APO
APO
APO
ATE
APO
AMO
AMO
APO
AMO
AMO
MC
AMO
AMO
AMO
AMO
AMO
APO
ATE
ATE
AMO
APO
AMO
AMO
APO
ATE
APO
ATE
AMO
APO
APO
APO
ATE
APO
ATE
APO
APO
APO
APO
APO
AMO
V
Sk
Sq
Sl
V
S:
S
S:
Q
V
B
Sq
Sq
Sr
S
S
Q
S
Sq
X:
K
Sq
S
Cg
S
Xk
S:
Q
V
X
Sq
S
C
Q
S
Sq
Sq
Sk
Sq
S
X
Q
S
Sk
Sq
X
Q
O
L
Xk
X
Xk
V
S:
Q
S:
Sq
S:
Sq
S:
S:
S:
S:
−0.2417
0.2964
0.1664
0.7720
−0.1423
−0.3718
−0.1185
−0.2466
−0.2027
−0.4932
0.0687
−0.0574
−0.0306
0.0391
0.1462
0.5354
−0.1596
−0.1655
−0.1338
−0.0074
−0.1459
0.2182
0.1124
0.2861
0.3304
0.5267
−0.0370
0.3880
0.2164
−0.1323
−0.5133
0.2284
−0.2762
−0.2085
−0.3287
−0.2154
−0.1734
−0.1015
−0.1632
−0.1613
0.0809
−0.0277
0.0296
−0.0697
−0.0165
−0.0335
−0.0201
−0.0597
0.0264
0.0481
−0.0064
0.5027
−0.0829
0.3277
0.0442
0.5743
0.2799
0.0223
−0.2265
−0.2690
0.1003
−0.2689
0.0553
−0.1584
0.0664
−0.0018
−0.1112
−0.0505
−0.1780
0.0140
−0.0573
0.5135
0.1469
0.5702
0.0574
−0.1000
0.4297
0.1758
0.2730
0.2899
0.2502
0.5264
0.1050
−0.0917
0.5336
0.3281
0.1494
0.3114
−0.0330
−0.0748
0.6373
0.3880
0.3049
0.2553
−0.1613
−0.3559
−0.5767
0.2743
−0.1628
−0.1543
0.2430
−0.2316
−0.1468
−0.2042
−0.2093
−0.1230
−0.0966
−0.1113
0.2749
−0.1754
−0.1871
−0.2121
−0.1547
0.1296
−0.1034
−0.2909
0.0301
0.1004
0.2639
0.2072
−0.5764
0.0138
−0.0557
−0.0380
−0.0459
−0.0136
0.0267
−0.0248
0.0080
−0.0196
−0.1006
−0.1395
−0.0468
0.0322
0.0732
−0.0144
−0.0714
−0.1680
0.0103
0.0425
−0.0099
0.0817
−0.0086
−0.1426
0.0131
−0.0039
0.1665
0.0135
−0.1665
0.0143
0.1182
−0.2621
−0.1422
0.1612
−0.4228
0.0349
17.0
18.5
16.8
18.4
17.3
15.5
15.4
20.3
17.8
19.6
18.5
19.1
18.5
21.9
18.4
20.5
18.1
18.9
22.5
14.5
17.2
19.5
17.0
19.6
19.0
19.9
16.6
16.0
19.8
20.0
19.3
14.6
19.5
19.4
17.2
16.8
18.5
21.1
18.5
19.1
16.8
21.0
22.9
22.4
17.7
17.9
16.2
18.2
20.4
18.2
19.2
19.3
20.4
24.8
16.2
15.5
19.2
19.5
17.4
16.8
16.3
18.7
19.2
T
Reference
4.72
∗
6.66
∗
5.10
∗
4.74
3
4.46
∗
6.49
∗
3.61
∗
6.37
∗
3.86
3
4.91
∗
3.20
∗
2.90
∗
3.84
3
3.65
∗
4.05
∗
5.70
5
5.14
∗
7.56
∗
4.19
5
6.98
∗
6.50
3
4.59
3
3.02
3
5.31
3
4.51
3
6.00
3
8.50
∗
3.90
∗
3.26
3
2.84
3
3.99
3
3.37
3
4.01
3
3.13
3
3.74
3
3.13
3
4.93
3
3.22
3
4.44
3
4.82
3
6.91
3
6.87
3
5.41
3
5.39
3
3.40
3
3.11
∗
4.58
∗
7.72
5
5.47
∗
6.12
∗
3.09
∗
5.71
∗
4.41
∗
3.19
∗
6.25
∗
3.10
∗
7.43
∗
5.31
∗
3.66
∗
3.09
∗
4.40
∗
3.33
∗
2.94
∗
(continued on next page)
266
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Table 1 (continued)
Number and name
Provisional designation
Orbit
Type
2000 JG5
2000 JQ66
2000 KL33
2000 MU1
2000 NM
2000 OJ8
2000 PG3
2000 RW37
2000 SY162
2000 WC67
2000 WF6
2000 WJ10
2000 WJ63
2000 WK10
2000 WL10
2000 WL63
2000 WM63
2000 WO107
2000 XL44
2000 YA
2000 YF29
2000 YH66
2000 YO29
2001 AE2
2001 CC21
2001 DU8
2001 EB
2001 EC
2001 FY
2001 HA8
2001 HK31
2001 HW15
2001 JM1
2001 JV1
2001 MF1
2001 OE84
2001 PD1
2001 QA143
2001 QQ142
2001 SG10
2001 SG286
2001 SJ262
2001 SK162
2001 TC45
2001 TX16
2001 TY44
2001 UA5
2001 UC5
2001 UU92
2001 UY4
2001 VG5
2001 VS78
2001 WA25
2001 WG2
2001 WH2
2001 WL15
2001 XN254
2001 XR1
2001 XS1
2001 XS30
2001 XU30
2001 XY10
2001 YE1
APO
AMO
AMO
APO
APO
AMO
APO
APO
AMO
AMO
AMO
AMO
AMO
APO
APO
AMO
APO
ATE
AMO
APO
APO
APO
APO
AMO
APO
AMO
AMO
APO
AMO
AMO
AMO
AMO
AMO
APO
AMO
AMO
AMO
AMO
APO
APO
APO
AMO
AMO
APO
MC
AMO
APO
AMO
AMO
APO
APO
AMO
APO
APO
AMO
AMO
AMO
APO
AMO
APO
APO
ATE
APO
S:
R
S
S
Sr
Sr
D
C
Sq:
X
Sq
Xk
Sq
X
Xc
S
S
X
S
Sk
S
Xk
C
T
L
S
Sl
Sq
S
C:
X
X
S
Sq
Sk
S
K:
Sk
Sq
X
D
C:
T
Sq
X
X
Sq
X
T
X
Sq
S
S
Sk
X
Sk
S
Sq
Cb
Xc
Sq
Sk
T
Slope
0.1709
0.4014
0.3325
0.2650
0.2645
0.8912
0.1129
0.1228
0.3689
0.1118
0.3045
0.0740
0.3869
0.2051
0.3429
0.5154
0.2668
0.3619
0.2565
0.5054
0.3397
0.0354
0.5668
0.4807
0.5482
0.5774
0.1184
0.4578
−0.1543
0.3964
0.1685
0.4320
0.1545
0.2617
0.3952
0.4288
0.2827
0.1837
0.2415
1.2455
−0.0494
0.5451
0.0674
0.3294
0.3215
0.2092
0.2903
0.5228
0.1920
0.2176
0.3394
0.2973
0.2598
0.2472
0.2623
0.4937
0.1650
0.1222
0.0369
0.1909
0.2047
0.6117
PC2
−0.5579
−0.2203
−0.1983
−0.2776
−0.2812
0.3498
0.1902
−0.0981
0.4058
−0.1327
0.1022
−0.1820
0.4195
0.2908
0.0344
−0.0725
0.3169
−0.0498
−0.0491
−0.0409
0.2441
0.2222
0.2436
0.0529
−0.0524
−0.0567
−0.1831
−0.1350
0.4436
0.2757
0.2479
−0.1749
−0.1214
−0.1227
−0.1748
0.0441
−0.0403
0.0108
0.3184
0.3639
0.2849
0.2713
0.0230
0.2983
0.4383
−0.0881
0.4534
0.3318
0.2353
−0.1558
−0.2435
−0.2064
−0.0519
0.3780
−0.0668
−0.1993
−0.1272
0.3537
0.1116
−0.0616
0.0281
0.2376
PC3
H Mag
−0.0166
−0.0499
−0.0468
−0.0259
−0.0603
0.0339
0.0312
−0.1457
−0.0624
−0.0243
−0.0856
−0.0577
0.0015
−0.0089
−0.0231
−0.0570
−0.0421
0.0778
−0.0387
0.0067
−0.1048
−0.0539
−0.1098
−0.0668
−0.0183
0.0056
0.0232
−0.0223
0.0146
−0.0072
0.0028
−0.0220
0.0050
0.0252
−0.0257
0.0519
−0.0250
0.0629
0.0318
−0.0237
−0.0085
0.0104
0.1085
0.0137
−0.0058
0.0642
0.0275
−0.1421
0.0021
−0.0288
−0.0211
−0.0618
−0.0570
−0.0078
−0.0806
0.0590
0.0019
0.0172
−0.0127
−0.0501
−0.0441
−0.0799
18.3
18.1
19.7
19.9
15.6
16.8
16.2
20.2
19.3
19.1
18.7
20.6
20.9
18.5
18.0
20.4
20.2
19.4
18.0
23.6
20.2
17.5
18.0
19.2
18.4
16.4
17.3
18.6
18.9
16.9
21.0
20.2
19.0
21.3
16.8
17.8
18.2
19.6
18.5
20.3
21.1
19.6
17.9
19.1
14.1
20.3
17.4
21.3
19.8
18.4
16.7
15.5
18.7
16.3
20.0
18.3
17.5
17.4
18.8
17.5
19.9
20.4
20.6
T
Reference
4.40
∗
3.56
∗
3.60
∗
4.71
∗
2.93
∗
3.31
∗
2.55
∗
5.09
∗
3.44
∗
3.09
∗
3.04
∗
3.59
∗
3.02
∗
4.25
∗
2.72
∗
4.58
∗
5.87
∗
6.23
∗
3.50
∗
3.21
∗
4.47
5
5.04
∗
3.36
∗
4.87
5
5.91
5
3.85
∗
4.07
∗
2.91
∗
3.89
∗
3.31
∗
3.25
∗
4.40
∗
4.52
∗
4.08
∗
3.03
∗
3.43
∗
3.49
∗
3.44
∗
4.64
∗
4.54
5
4.77
5
2.98
∗
3.77
5
3.31
∗
2.77
∗
3.35
∗
3.94
∗
2.87
∗
2.80
∗
4.23
∗
3.28
∗
3.94
∗
3.90
∗
3.56
∗
3.65
∗
3.56
∗
3.35
∗
4.95
∗
3.14
∗
4.93
∗
3.33
∗
6.66
∗
3.82
∗
(continued on next page)
Spectral properties of near-Earth objects
267
Table 1 (continued)
Number and name
Provisional designation
Orbit
Type
Slope
PC2
PC3
H Mag
T
Reference
2001 YK4
2002 AA
2002 AD9
2002 AH29
2002 AK14
2002 AL14
2002 AL31
2002 AQ2
2002 AT4
2002 AU5
2002 AV
2002 BA1
2002 BK25
2002 BM26
2002 BP26
2002 CS11
2002 CT46
2002 DH2
2002 DO3
2002 DQ3
2002 DU3
2002 DY3
2002 EA
2002 EC
APO
APO
APO
AMO
APO
ATE
APO
AMO
AMO
APO
APO
AMO
APO
APO
AMO
AMO
AMO
APO
APO
AMO
APO
AMO
APO
AMO
X:
Sq
L
K
V:
Ld
X
S
D
X
K
S
Sk
X
X
X:
Sr
Ch
X:
Sq
Sq
Xk
L
X:
0.2210
0.2389
0.6439
0.3581
−0.0590
0.9826
0.2859
0.3978
0.8578
0.3663
0.2958
0.4377
0.1673
0.4393
0.2082
0.2506
0.3123
0.0228
0.2489
0.0135
0.2518
0.2668
0.4936
0.2545
0.2629
−0.1375
0.0097
0.0702
−0.4693
0.1559
0.2359
−0.0367
0.2713
0.2713
0.0557
−0.2128
0.0427
0.2702
0.2083
0.2713
−0.3365
0.3159
0.5103
0.0272
−0.2181
0.1899
0.0218
0.3439
0.0195
0.0363
−0.0828
−0.0019
0.0365
0.0825
−0.0202
−0.1564
0.0104
0.0104
0.0109
0.0042
−0.0306
0.0408
0.0498
0.0104
0.0107
0.0128
0.0566
−0.0132
0.0157
0.0928
−0.0345
−0.2688
18.5
19.5
16.5
21.7
21.7
17.8
24.4
18.6
20.9
17.7
20.7
21.7
18.1
20.1
19.3
21.6
20.9
20.3
22.0
23.8
20.8
18.6
22.4
23.3
2.83
5.42
3.58
3.29
5.97
6.26
5.34
3.08
3.95
3.56
3.13
3.73
3.13
3.86
3.97
3.70
3.30
3.52
3.83
4.74
5.44
4.43
4.69
3.51
∗
∗
∗
∗
∗
∗
5
∗
5
∗
∗
∗
∗
∗
∗
∗
∗
∗
∗
5
5
∗
∗
∗
∗—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).
presented here, where these objects are denoted by “*” in the
Reference column. Observing circumstances for all newly
reported measurements are presented in Appendix A and the
newly reported spectra appear in Appendix B and are available in digital format at http://smass.mit.edu/.
For SMASS observations prior to 1997, visible wavelength measurements were made almost exclusively with the
Mark III long-slit spectrograph coupled to the MichiganDartmouth-MIT Observatory (MDM) 2.4-m telescope located on the southwest ridge of Kitt Peak, Arizona. Our
observing procedures carried out there (fully described in
Xu et al., 1995; Bus, 1999; Bus and Binzel, 2002a) were
followed closely at all sites. Particular details in common
include using a long slit oriented north–south to simultaneously image the object spectrum and the background sky,
with a slit width several times wider than the seeing disk for
the best possible photometric precision. Comparable spectral
resolution, typically λ/λ ∼ 100 was obtained with all telescope/instrument/slit/detector configurations. Spectral exposures, typically 900 s or shorter (to minimize the accumulation of cosmic-ray hits on the detector) were nearly always
made while the object was within one hour of the meridian to minimize any effects of atmospheric dispersion. For
all observations we utilized Hyades 64 and 16 Cyg B as our
primary reference stars for the solar analog spectrum. For
further sky coverage, we also utilized solar analog reference
stars selected from Landolt (1973) that were verified to be
within 1% of our primary reference stars, as detailed in Bus
and Binzel (2002a).
In 1998 a new collaboration began providing to the
SMASS program visible wavelength measurements made
using the Double Spectrograph on the Palomar Observatory 5-m Hale telescope (co-author AWH, Palomar Principal
Investigator). Our use of this telescope and instrument combination is fully described in Binzel et al. (2001a). SMASS
NEO visible wavelength measurements were also begun using the RCSP spectrograph on the Kitt Peak National Observatory 4-m Mayall telescope (Binzel, P.I.) in 1999, as
detailed in Binzel et al. (2001b). While our CCD spectra obtained at MDM, Palomar, and Kitt Peak generally ranged
only out to 0.92-µm, a program to extend measurements out
to 1.6-µm was begun at the NASA Infrared Telescope Facility (IRTF) in 1997. This so-called “SMASSIR” survey
utilized a low-resolution “asteroid grism” system designed
by one of us (RPB) and described in Fig. 2 of Binzel et
al. (2001a). The performance of the asteroid grism and the
methods utilized for the acquisition, reduction, and calibration of these data are detailed in Burbine (2000) and Burbine
and Binzel (2002).
In this paper, we also present results from a one night
visible spectroscopy run on the Keck II 10-m. This Keck
telescope night was allocated to obtain spectral measurements of 4660 Nereus in support of the NASA collaboration with the MUSES-C mission. (Nereus was the initially
planned MUSES-C target.) We utilized LRIS (low resolution imaging spectrometer) to obtain spectra over the wavelength range of 0.5- to 1.0-µm with a 300 line/mm grating.
The basic observing, reduction, and solar calibration procedures described above were also applied to these Keck
measurements. While Nereus was not available during the
entire night, four other NEOs were measured as targets
of opportunity. Finally, we also incorporate results from
268
R.P. Binzel et al. / Icarus 170 (2004) 259–294
the newly operational Magellan I 6.5-m telescope at Las
Campanas, Chile. There we utilized the Boller and Chivens
spectrograph, also using a 300 line/mm grating to cover the
wavelength range of 0.5- to 0.9-µm. Similarly, our standard
observing, reduction, and solar calibration procedures were
followed.
3. Taxonomic analysis
3.1. Taxonomic classification
Our first analysis step is to place the newly observed objects into the taxonomic system of Bus (1999). The Bus
taxonomy is based on a uniform data set of 1341 mainbelt asteroid spectra from the SMASSII survey measured
with the 2.4-m telescope and Mark III spectrograph at the
MDM Observatory (Bus, 1999; Bus and Binzel, 2002a). As
our NEO survey utilized a wide variety of telescopes and
spectrograph combinations, we chose to re-observe a small
subset of main-belt objects for direct comparison and correlation with the SMASSII survey results. To quantify our
comparison, we calculate spectral slope values and principal component scores using the Bus (1999) slope definition
and eigenvectors. Sufficient data were obtained to compare
the principal component scores for 19 SMASSII main-belt
asteroids re-observed at Palomar and 24 re-observed at Kitt
Peak. As the Kitt Peak data generally do not extend below
0.50-µm, extrapolation was necessary to the 0.44-µm lower
limit of the eigenvectors. This extrapolation was made following the curvature (or linearity) occurring in the 0.5- to
0.6-µm spectral range. Both Palomar and Kitt Peak sets show
a slight but mutually consistent offset from the SMASSII
slope and PCA scores. This is a purely empirical offset, perhaps due to slightly different spectrophotometric responses
of the various systems. As shown in Fig. 1, this offset is too
small to affect any taxonomic classification except in the
immediate vicinity of the class boundaries where a natural
ambiguity always exists (Bus, 1999). However because we
seek to perform statistical trend analysis, we must account
for this offset in order to achieve a data set that is as internally consistent as possible. Letting Slope, PC2 , PC3 be the
values on the system of Bus (1999) and letting Slope0 , PC20 ,
PC30 be the initially calculated values for the new Palomar
and Kitt Peak data reported here, the transformations to the
SMASSII 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 presented
in Table 1. We note that no transformations were applied
to the 106 NEOs observed during the SMASSII survey using the MDM 2.4-m telescope as these are formally on the
SMASSII system. (These SMASSII measurements are distinguished in Appendix A as having been made using MDM
2.4-m telescope.) Similarly, 10 NEOs observed in SMASSI
(Xu et al., 1995, also using the MDM 2.4-m telescope)
have their principal component scores recalculated with the
SMASSII eigenvectors but with no transformation applied.
The resulting distribution of Slope and Principal Component
scores comparing the near-Earth and main-belt populations
are plotted in Fig. 1.
Fig. 1. Principal component space within the Bus (1999) taxonomy system designed for visible wavelength CCD spectra. The 1341 SMASS main-belt asteroids
defining this system are depicted by their classification letters while 400 SMASS near-Earth and Mars-crossing objects are depicted in red. NEOs show a greater
dispersion in spectral properties, and most notably, span the once empty gap between the S- and Q-complexes (Binzel et al., 1996, 2001a). The green line
denotes the defined boundary between the C- and X-complexes, where objects close to this line have a natural ambiguity in their taxonomic assignments.
D-type objects in the upper right are candidates for extinct comets. The horizontal line at lower right represents the magnitude of the average transformation of
Slope values for Palomar and Kitt Peak measurements to the Bus system. The vertical downward arrow is the PC2 transformation (a constant). The magnitude
of these two lines is also representative of the typical uncertainty for placement of any object within principal component space.
Spectral properties of near-Earth objects
Assignment of a feature-based taxonomic classification
to each object was made following the description of Bus
and Binzel (2002b) as originally developed by Bus (1999).
All taxonomic assignments are presented in Table 1. Several
of the NEOs observed in SMASSI (Xu et al., 1995) have
their taxonomic types updated to the Bus system, but all remain in the same “complex.” (For example 2074 Shoemaker
is revised from “S” to “Sa” but remains in the S-complex.)
For some objects, only SMASSIR data are available over
the range 0.9- to 1.65-µm. Since the Bus taxonomic classes
are defined over visible wavelengths below this range, it is
not formally possible to place these near-infrared only measurements into an existing taxonomic system. However, the
spectral characteristics over the 0.9- to 1.65-µm range are
generally recognizable as being consistent with S, C, or X
classes. Within Table 1, we list these SMASSIR-data-only
results as S:, C:, X:, where the colon denotes the uncertainty
of the taxonomic assignment.
Taxonomic assignments have natural ambiguity near
class boundaries, particularly for cases of low signal-tonoise ratio (SNR) spectra. Cases of completely ambiguous
and noisy spectra have been deleted from the data set. These
cases comprise only about 1% of the total presented here.
The boundary between the C- and X-complexes gives rise
to the greatest natural ambiguity for both high quality and
less than perfect data. Similarly, some ambiguity can exist
between the X- and S-complexes based on the quality of the
spectral data for revealing the presence or absence of an absorption band beyond 0.8-µm, a characteristic reflected by
principal component PC2 . To resolve potentially ambiguous cases, our taxonomic assignments have been made using
both the principal component scores and the best match of
the spectra to the defined ranges for each class provided by
Bus and Binzel (2002b). Asteroid 2100 Ra-Shalom provides
an example where high SNR data can prove ambiguous as
well. The principal component scores and spectral characteristics of Ra-Shalom place it at the boundary between Xcand C-types, where Bus and Binzel (2002b) denote it as Xc.
With the addition of near-infrared data, the continued decreasing slope is more characteristic of C-type than X-type
objects. For this tabulation, we place Ra-Shalom in the Ccategory.
Five objects have sufficiently unusual or relatively low
SNR spectra that place them outside the range of the Bus
classes or make their taxonomic assignment fully ambiguous. For one of these, 3908 Nyx, we follow the analysis of
Burbine (2000) and place it in the “V” class. For the other
four we choose to maintain the designation “U” as originally
listed by Bus and Binzel (2002b) and Binzel et al. (2001a).
(5646) 1990 TR is ambiguous between Q and V. (7888) 1993
UC may be an extreme form of an A-type, but falls well outside the range for this class. (8566) 1996 EN appears to be
an extreme form of V-type, but falls well outside the range
for this class. (20043) 1993 EM is presented in Binzel et al.
(2001a) and displays an ambiguous low SNR neutral spectrum.
269
Objects observed on multiple nights or telescopes (as
noted in Appendix A) have their final results based on
a weighted average of all measurements. For three cases,
datasets having significantly higher signal-to-noise (SNR)
are preferentially used for the tabulation and analysis presented here. 4660 Nereus is recognized to be an Xe-type
based on high SNR favorable apparition measurements obtained at Kitt Peak and Palomar, as discussed in Binzel et
al. (2004a). Keck measurements for Nereus show a spectral slope consistent with this classification, but have a lower
SNR from a faint (V20.5) apparition. For (5587) 1990 SB,
KPNO 4-m spectra reveal an Sq type. Lower SNR SMASSII
measurements are consistent with this result. For 1994 AW1,
lower SNR SMASSII data are consistent with both Sa and L,
with Sa tabulated in Bus and Binzel (2002b). Higher SNR
data subsequently obtained at Palomar are more consistent
with an L-type classification that we consider a more secure
result, which we tabulate here. As an additional note, the
Mars crosser 1011 Laodamia is reported to be an Sr-type by
Bus and Binzel (2002b). A higher SNR spectrum obtained
using the KPNO 4m is more consistent with an S-type designation, which we tabulate here.
3.2. Observed taxonomic distribution
Figure 2 reveals that the observed taxonomic distribution
of the NEO population effectively spans the full breadth of
the main-belt diversity, with 25 of the 26 Bus taxonomy
classes represented within the SMASS NEO sample. Only
the Bus class Cgh, typically found among outer main-belt asteroids (beyond 2.7 AU) is not recognized within the current
SMASS NEO sample. A summary of the taxonomic classification statistics is presented in Table 2.
Within Table 2, we also seek to collate results among major groups that are consistent with what are referred to as
“complexes” by Bus (1999). We make these “complex” assignments for the purpose of establishing taxonomy input
parameters to the bias-corrected population analysis performed by Stuart (2003) and Stuart and Binzel (2004). We
differ slightly from the Bus (1999) use of the term “complex” in two ways. The first is our grouping the “Sq” and
“Q” designations into a complex we denote as “Q.” We
make this grouping because these types of objects are common in the near-Earth population but are rare or absent in
the Bus (1999) main-belt sample. The second is our treatment of the degeneracy of the X-class. As described by
Tholen (1984), the designation “X” denotes spectrally similar objects that are best distinguished by their albedos. From
highest to lowest albedos, the distinct classes are labeled
as E, M, P. Two factors allow us to address this degeneracy. The first is the availability of albedo data for a sample
of NEOs (Delbo et al., 2003). The second is the correlation emerging between the E- and Xe-classes. Within the
Bus (1999) system, a 0.49-µm feature that distinguishes the
Xe-class appears fully consistent with the high albedo Eclass, a result exemplified by 4660 Nereus (Binzel et al.,
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R.P. Binzel et al. / Icarus 170 (2004) 259–294
Fig. 2. Diversity of taxonomic classes observed among the near-Earth object population. Fully 25 of the 26 Bus (1999) taxonomic classes are identified in the
vicinity of the Earth, suggesting a broad contribution to the population from the main-belt.
Table 2
Summary of NEO taxonomic classification statistics
Bus class
A
B
C
Cb
Cg
Ch
Cgh
D
K
L
Ld
O
Q
R
S
Sa
Sk
Sl
Sq
Sr
T
V
X
Xc
Xe
Xk
U
SMASS sample
1 (4)
5 (3)
13 (2)
3 (2)
1
1 (3)
4
7
7 (2)
2 (1)
6
18 (2)
1
76 (40)
2 (6)
13 (4)
6 (2)
62 (9)
12 (2)
5 (1)
14
31 (2)
6 (1)
2 (4)
9
3 (1)
Complex
A [A]
C [B, C, Cb, Cg, Ch, Cgh]
D [D, T]
E [E, Xe]
M [M]
P [P]
O [O]
Q [Q, Sq]
R [R]
S [S, Sa, Sk, Sl, Sr, K, L, Ld]
V [V]
X [X, Xc, Xk]
U [U]
SMASS sample
1 (4)
23 (10)
9 (1)
3 (4)
4
6
80 (11)
1
125 (57)
14
41 (3)
3 (1)
Other dataa
15 (1)
1
3
19
20 (6)
2
6
Total
1 (4)
38 (11)
9 (1)
4 (4)
3
4
6
99 (11)
1
145 (63)
16
47 (3)
3 (1)
Numbers in parentheses are a separate tabulation for the Mars-crossing (MC) population. Brackets [ ] depict the taxonomic classes used for the consolidation
into each complex.
a Statistics within this column are based on additional measurements of NEOs, representing the dedicated work of many observers, including McFadden
et al. (1985), Cruikshank et al. (1991), Hammergren (1998), Hicks (1998), Rabinowitz (1998), Whiteley (2001), Angeli and Lazzaro (2002). A tabulation of
most of these results appears in Binzel et al. (2002), for which updates and references to the original sources are maintained at http://earn.dlr.de/nea/.
2004a). Albedo information (Harris and Lagerros, 2002;
Delbo et al., 2003) is available that allows the degeneracy
to be resolved for three objects classified as “X” in Table 1:
5751 Zao is likely an E-type while 1999 JM8 and 2000
BG19 are likely P-types. Similarly, 3691 Bede listed as an
Xc-type is distinguishable by its albedo (Delbo et al., 2003)
Spectral properties of near-Earth objects
and is likely a P-type. For these four objects plus two Xe
objects we tabulate separate entries among the E, M, and P
“complexes” in the right-hand columns of Table 2. An interesting outcome of our trying to resolve the degeneracy
within the X-complex is to estimate what percentage of these
observed objects might be considered “dark” and “bright”
along the lines of the analysis performed by Morbidelli et
al. (2002b). Within the SMASS sample, a total of 22 objects
have sufficient albedo information or diagnostic spectral information to be classified distinctly as E, M, P, Xc, Xe, or
Xk. Of these, 10 out of 22 or 45% fall within the P- and Xcclasses that might be considered as “dark” objects, while the
other 55% might be broadly considered as being “bright.”
Reducing the taxonomic classifications into complexes
also facilitates a comparison between the SMASS results
and other available published data on NEOs. (Complexes effectively represent a “least common denominator” between
multiple taxonomic systems.) Results from other sources are
available for more than 70 objects for which no SMASS
observations have been obtained. (See reference footnote
in the table as these additional data represent the work of
many dedicated observers.) Totaled together, taxonomic information is currently available for more than 370 NEOs and
nearly 100 Mars-crossers. SMASS and all other programs
concur in that the S-, Q-, C-, and X-complexes account for
90% of all objects. SMASS differs in having a higher X:C
ratio than that measured by other programs. This may simply be a selection effect where relatively neutral colors from
filter measurements are most easily branded as “C-types.”
Within SMASS, the full visible wavelength coverage of
CCD spectra may better enable an X:C distinction, although
this is still subject to the X:C ambiguities described above.
A comparison between the observed near-Earth and
Mars-crossing populations is displayed in Fig. 3. While
roughly 2/3 of all observed NEOs fall into either the S(40%) or Q- (25%) complexes, the S-complex (65%) alone
dominates the Mars-crossers. The low proportion of Q-
271
complex (10%) objects among Mars-crossers may be a size
effect (Section 5), with the closer proximity of NEOs allowing their discovered and spectrally measured population to have smaller average sizes. The greater abundance
of D-types among NEOs may also be a size/albedo selection effect, although the apparent preferential origin of Dtypes from Jupiter-family comets (Section 4) and their rapid
dynamical evolution with very little time spent as Marscrossers may be an important factor.
Finally, we can compare our observed taxonomic distributions with the modeling assumptions of Morbidelli et
al. (2002b) who assumed the observed ratio of “dark” to
“bright” objects to be 0.165 for H < 20. Within the total
SMASS sample (Table 2), we count 207 objects having H <
20 in the A-, E-, M-, O-, Q-, R-, S-, V-, and U-complexes (as
well as Xe- and Xk-class objects) as being “bright.” Among
the C-, D-, and P-complexes (as well as the Xc-class), we
count 30 “dark” objects. The 19 objects tabulated in the Xcomplex are problematic. We assume that 45% (9 objects)
of these are “dark,” where this percentage comes from the
“dark/bright” assessment of the X-complex discussed above.
Combining, we find the SMASS observed dark/bright ratio among H < 20 objects is (30 + 9)/(207 + 10) = 0.18,
a slightly higher ratio compared with the 0.165 value assumed by Morbidelli et al. (2002b). The greatest uncertainty
(assuming all observational errors are random) in our 0.18
value is our treatment of the X-complex. For example, if we
consider 9 ± 3 as a reasonable uncertainty for how many
X-complex objects to assign to the “dark side,” the uncertainty in our ratio becomes 0.18 ± 0.02, a range encompassing therefore compatible with the 0.165 assumed value of
Morbidelli et al. For H 20, the SMASS observations we
calculate a dark/bright ratio following the same procedures
as above. The resulting calculation yields (8 + 5)/(34 + 7) =
0.32 ± 0.06 as the observed dark/bright ratio for objects having H 20. (Here we assign 5 out of 12 X-complex objects
to the “dark side” and assume the uncertainty in this number
Fig. 3. Comparison of the observed NEO and Mars-crossing populations, grouped into broad taxonomic complexes. Nearly 90% of the observed objects fall
within the broad S-, Q-, X-, and C-complexes. S and Q dominate the NEOs, with 40 and 25%, respectively, while S dominates (65%) the MCs. V-type objects
represent the next largest category (4%) of observed NEOs, but appear absent among MCs. Relatively rare A- and E-complex objects are shown to be more
common among Mars-crossers, each comprising 4% of the observed MC population.
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R.P. Binzel et al. / Icarus 170 (2004) 259–294
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.g.,
∗, #, &, +) are tagged to select objects for visually correlating them from eccentricity to inclination space. The C ∗ symbols show a grouping of C-type objects
near a, e, i = 2.76, 0.42, 29, which may be a family of Mar-crossing objects. They derive from a relatively sparsely populated high inclination region of the
central 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 Flora itself; 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 concentration
of similarly rare types). If they come from the Flora region, their diffusion follows the route predicted by Morbidelli and Nesvorny (1999) in which slow
changes 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 followed the
same predicted diffusion paths from the Hungaria region. (The Hungaria region is rich in E-types which may be related to enstatite achondrite meteorites
Gaffey et al., 1992.) All tagged groupings are summarized in Table 3.
to be 5 ± 2.) Given the small number statistics and the uncertainties involved in making “dark” or “bright” assignments
from taxonomy alone, we believe it is an open question as
to whether there is a distinct increase in the number of dark
objects for H 20.
4. Source region analysis
4.1. Orbital distribution
While transfer from the main belt to the vicinity of the
Earth is a chaotic process (Wisdom, 1985), we investigate
whether the combination of taxonomic classes plus orbital
elements retains any signatures for tracing the origins of
NEOs. In particular, the slow diffusion of objects into Marscrossing orbits (Morbidelli and Nesvorny, 1999) may pro-
vide an observable trace. Figure 4 shows the NEO and MC
populations in semi-major axis (a), eccentricity (e), and inclination (i) space. The alphabet soup for the NEO population seen in Figs. 4A and 4B attests to their dynamical
mixing. However these panels do show some distinct signatures remaining within the MC population, possibly consistent with slow diffusion. In Table 3 we summarize our
findings for three possible groupings and discuss each group
in turn below.
Figures 4A and 4B show a cluster of five C and C-subtype
asteroids (all denoted by C ∗ ) located near a, e, i = 2.76,
0.42, 29. The dynamical evolution of these objects is likely
related to the adjacent 5:2 resonance. Whether or not this
grouping constitutes a “family” (implying a collisional origin from a common parent body) remains an open question
as these are common main-belt classes in this region. What’s
more, they are more diverse (comprised by C-, Ch-, and B-
Spectral properties of near-Earth objects
273
Table 3
Mars-crossing asteroid groups
Group
a
e
i
Members
Types
Potential source
1
2
1.94
2.20
0.06
0.25
22
5
E
A, Sa
Hungaria region
Flora region
3
2.76
0.42
29
2035 Stearns, 6249 Jennifer
2423 Ibarruri, 3858 Dorchester,
3920 Aubignan, 5275 Zdislava
3581 Alvarez, 6500 Kodaira,
5349 Paulharris, 5585 Parks,
5870 Baltimore
C and C-subtypes
No identified source, but a common type
for outer main belt
types) than is typical for a homogeneous family. Unlike the
other groupings in Table 3, this one has no apparent association with any previously known cluster.
Figures 4C and 4D focus on the inner belt and taxonomic
classes (A, E, Sa, V) that are generally interpreted to indicate
objects that have undergone a moderate degree of heating
and differentiation. Being rare taxonomic types, they offer
the opportunity for some traceability in their dynamical evolution. According to the diffusion model of Morbidelli and
Nesvorny (1999), evolution in eccentricity occurs first with
little or no change in semi-major axis until Mars-crossing
orbits (with the potential for Mars encounters) are reached.
We find two cases of rare taxonomic types that may be traceable across the Mars-crossing boundary as direct evidence of
this predicted diffusion. In the first case, we find that E-type
asteroids (often related to enstatite achondrite meteorites;
Gaffey et al., 1992) are abundant in the Hungaria region
located near a, e, i = 1.94, 0.06, 22. Two E-type objects apparently caught in the act of diffusing from the Hungarias are
2035 Stearns and 6249 Jennifer (denoted by the E& symbol
in Figs. 4C and 4D). A third object 2449 Kenos (denoted by
E+) matches in eccentricity, but not inclination. If related,
it has already begun a direct interaction with Mars affecting
its orbit plane. Gaffey et al. (1992) make a strong case for
E-type 3103 Eger to be derived from the Hungaria region,
where its evolved orbital elements (a, e, i = 1.40, 0.35, 20)
place it completely out of the area of Figs. 4C and 4D, requiring a much more evolved dynamical history than the objects
we list in Table 3.
A second case for diffusion in orbital eccentricity to become Mars-crossing objects is also found in the A and Sa
rich Flora region near a, e, i = 2.2, 0.16, 6. The visible spectra of A- and Sa-type objects suggest a mineralogy rich in
olivine, as does the Sr-class also seen in this region (Gaffey,
1984; Florczak et al., 1998) A group of four objects reside
just across the Mars-crossing boundary near 2.2 AU. In the
figure they are denoted by Sa# and A# and their identities
are given in Table 3. All four have the same inclinations as
Flora, making them prime candidates for having recently diffused from this region but not having yet begun substantial
interactions with Mars. Taken together, the apparent traces
of diffusion from the Hungaria region and the Flora region
provide strong observational support for the weak resonance
diffusion into the Mars-crossing population predicted by
Morbidelli and Nesvorny (1999).
Figure 4 (as well as Fig. 2 and Table 2) indicates that
V-types are rare or absent among observed Mars-crossers.
(16 observed V-types fall in the NEO category, none are observed in this sample among MCs.) This dichotomy implies
a low eccentricity main-belt origin for the V-types via the
ν6 and 3:1 resonances, a process in which objects pass very
rapidly (and therefore unlikely to be observed) through the
Mars-crossing phase in becoming NEOs (Morbidelli et al.,
2002a). High eccentricity objects can become NEOs via the
slow diffusion process of Morbidelli and Nesvorny (1999),
with a long residence time as Mars crossers. The rarity of
V-types among Mars crossers implies they do not have a
significant high eccentricity source. Most scenarios (e.g.,
Consolmagno and Drake, 1977; Binzel and Xu, 1993) predict that V-type NEOs (and HED meteorites) are derived
from Vesta or its apparently associated collision fragments.
The low orbital eccentricities of Vesta and the “Vestoids,” if
indeed they are the source of V-type NEOs, would produce
the predominance of V-type NEOs relative to Mars crossers
that is observed.
Finally we note a distinct variation in taxonomic distribution with respect to the jovian Tisserand parameter (T ),
as readily apparent in Fig. 5. We describe this more fully in
Section 4.3.
4.2. Taxonomic signatures from source regions
With the available data set of measured taxonomic properties for the NEO population, a marriage of observations
with theoretical modeling now becomes possible for a detailed examination of NEO sources. Bottke et al. (2002)
provide a model for calculating the probability of an object entering NEO space from one of five source regions:
the ν6 resonance, the Mars-crossing zone, the 3:1 resonance,
the outer belt, and Jupiter family comets. For each object in
our sample, Bottke (personal communication) provided the
model values for the five source regions based on the object’s
current orbital elements, where the sum of the five model
probabilities is equal to one. We couple these model probabilities with the NEO data consolidated into the “complexes”
of Table 2. We examine the resulting source region probability distribution for each complex separately. For example,
we take the Bottke model probability numbers for just the Ccomplex objects and sum the fractional probabilities within
each of the five source bins. The results for each of these
five source bins are normalized so that their total probabil-
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R.P. Binzel et al. / Icarus 170 (2004) 259–294
Fig. 5. Distinct differences are found between the T 3 and T > 3 populations when comparing their fractional abundances (number in class/total sample).
C-, D-, and X-complex objects (Table 2) are found to dominate for T 3, all of which are in low albedo categories (assuming the “X” objects represent the
P-class), consistent with the findings of Fernandez et al. (2001). These T 3 objects may be candidates for extinct comets (Wiessman et al., 1989, 2002).
ity is one. We follow this in turn for D-, E-, M-, P-, Q-, S-,
V-, and X-complexes, where Fig. 6A displays the results for
each complex. For the objects which we sampled, the ν6 resonance prevails as the source with the greatest contribution,
accounting for 46% of the source probability for the sum of
our entire sample (all complexes combined). Bottke et al.
(2002) (see their Table 3) predict a steady state NEO contribution of about 37% from the ν6 resonance. We account
for the greater ν6 dependence in our sample as being due to
observational selection effects: objects at the inner edge of
the asteroid belt are more easily discovered and more likely
to have their properties measured in a magnitude limited
survey. Source probabilities for our sample from the intermediate Mars-crossers (27%), the 3:1 resonance (19%), and
the outer belt (6%) match well with the values of Bottke
et al. (2002). Our sample has a lower (2% as compared to
6%) source probability from Jupiter family comets, an effect
likely due to the difficulty of discovery and observation of
high eccentricity and typically distant objects, especially if
they are dominated by classes having low albedos.
While differences in the source regions for one complex
compared to another are apparent in Fig. 6A, they are more
readily seen in Fig. 6B where we normalize each source region by the average for the total sample. As an example,
C-type objects contributed by the 3:1 resonance have a histogram value 0.25 in Fig. 6A, while the average contribution
from the 3:1 resonance for our entire sample is 19%. Figure 6B shows the normalized C-type contribution via the 3:1
resonance to be given as 1.31 (leftmost striped bar), where
1.31 is the quotient of 0.25 normalized (divided) by 0.19.
If all taxonomic types were equally likely to be contributed
from all sources, all histogram values of Fig. 6B would be
unity.
Distinct variations from unity are revealed in Fig. 6B
that represent the signatures of higher (or lower) than average contributions to the NEO population from the five
sources. The greater than unity values for the 3:1 and outer
belt sources within the C-complex reveal that these regions
deliver a proportionally larger fraction of C-type objects to
near-Earth space. Such a finding is fully consistent with the
predominance of C-type objects at the 3:1 resonance and beyond. Similarly, the dominance D-types being contributed
from Jupiter family comets is strikingly shown, with important implications for deriving the extinct comet fraction from
the NEO population, as discussed below. Because S-types
comprise the greatest fraction of our total sample, it can be
expected that their source contributions will closely approximate the average. The high degree of similarity seen between
the S- and Q-complexes may have an important implication:
if S- and Q-asteroids are related (such as Q-types just being
“fresh” S-type surfaces), then a necessary (but not sufficient)
condition is that they show similar structures for their source
regions. V-types also show a very similar profile relative to
the S-types. While the Fig. 6B histogram predicts relatively
equal entry of V-type NEOs via the Mars-crossing route and
resonance routes, the lack of observed V-type Mars crossers
(noted above) suggests a quick passage from a resonance origin through the Mars-crossing phase, into NEO space.
Finally, in our examination of source region signatures
we evaluate the E-, M-, P- and X-complexes. X-complex objects show a very strong contribution to the JFC population
that we interpret as arising from the fraction of low albedo
P-types, which by definition are part of the X-complex. We
note that only a total of 10 X-type NEOs have albedo information for discriminating them into the E, M, P classes
(Table 2). Thus the E, M, P signatures in Fig. 6B are based
on a very limited sample and the resulting implications can
only be considered preliminary. Among E-types, the higher
than average contribution from the innermost ν6 resonance
is consistent with the high abundance of E-types in the adjacent Hungaria region. Both M- and P-types show evidence
for high concentrations from the outer belt, where primitive
Spectral properties of near-Earth objects
275
Fig. 6. A. Contributions by taxonomic complexes (Table 2) to the NEO population are presented as a function of the source regions modeled by Bottke et al.
(2002). These regions: ν6 resonance, Mars-crossing (MC) zone, 3:1 resonance, outer belt (OB), and Jupiter-family comets (JFC), respectively, contribute 46,
27, 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 sources is
unity. The ν6 resonance is seen as the most prolific provider for all classes. B. Signatures of source region contributions for each taxonomic complex. These
signatures are revealed by rationing (e.g., by 0.46 for ν6) the contribution in each zone by its average and comparing the results with respect to unity (dashed
line). Many distinctive sources are revealed: C-types have a proportionally higher contribution from the 3:1 resonance and outer belt while D-types are strongly
sourced from Jupiter family comets. The sample size for E-, M-, P-types is small, thus their trends are only preliminary indications: E-types preferentially
enter via the ν6 resonance. M-types in the sample show an origin preference in the outer belt, perhaps suggesting they may be examples of “primitive” objects
within 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 provenance
from 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 expected to reflect the average. Both Q- and V-types also reflect the average, suggesting
similar main-belt region origins as S.
(not strongly heated) asteroids predominate. While P-types
are known to predominate in the outer belt, seeing an M-type
spike in the outer belt is confusing, but is based on a sample
of only three objects. We discuss the M-types in Section 6.
4.3. Tisserand parameter and comet fraction
In a classic analysis, Fernandez et al. (2001) found a dependence on albedos related to the Tisserand parameter T
(with respect to Jupiter):
a
aj
+ 2 cos(i) 1 − e2
,
T=
a
aj
where aj is the semi-major axis of Jupiter. Inner and outer
Solar System objects having T 3 were found to have distinctly lower albedos, which led Fernandez et al. to conclude
these objects were candidates for extinct comets. The value
of T 3 denotes objects that are dynamically coupled to
Jupiter, as is the case for Jupiter family comets. Objects
may change their Tisserand value through non-gravitational
forces (such as cometary outgassing), by interacting with
other planets, and as a result of the eccentricity of Jupiter.
Because NEOs do interact with inner planets, T 3 is not
a rigid boundary for objects that may have originated as
Jupiter family comets. Similarly, objects originating in the
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R.P. Binzel et al. / Icarus 170 (2004) 259–294
main-belt with T > 3 can undergo planetary encounters that
cast them into the control of Jupiter. Thus objects that originated as Jupiter family comets may now have T > 3, just
as objects originating in the main belt can now have T 3.
Cometary outcasts across this border may be recognized (or
suspected) based on their very low albedos or unusual taxonomic classes, as discussed below. Asteroidal outcasts might
most likely manifest themselves in T 3 orbits as S- and Qtype objects—as these most common types of NEOs would
have the greatest random chance of being scattered across
the T = 3 boundary.
Here we make a similar but independent analysis to that
of Fernandez et al. (2001), utilizing T 3 as a distinguishing parameter but focusing only on the near-Earth population. Rather than albedo, we use taxonomic classes as
our variable. Current discovery statistics (IAU Minor Planet
Center) show that 7% of all catalogued NEOs have T 3.
Within the sample population of SMASS, 6% of the observed objects have T 3, a close representation for the discovered population. While C-, D-, Q-, S-, and X-complexes
have more than one T 3 member in our sample, our T 3
sample of NEOs is dominated (50%) by C-, D-, and X-class
objects as shown in Fig. 5. For our analysis of the T 3
population, we make two assumptions: First, the T 3 Sand Q-types within our sample are randomly scattered mainbelt objects as discussed above (where S and Q are the most
likely letters to be drawn from the main-belt Scrabble® bag).
Second, we assume the X-type objects having T 3 are actually P-types based on their very strong signatures of being
derived from the outer belt (Fig. 6B). (Of the 6 X-complex
NEOs within our T 3 sample, we have albedo information for only one, 1999 JM8. From radar measurements by
Benner et al. (2002), 1999 JM8 is inferred to have a low
visible wavelength albedo.) This predominance of the C, D,
and (assumed) P objects for T 3 is consistent with the
low albedo correlation shown by Fernandez et al. (2001).
Starting with 7% of the discovered NEO population having T 3, and our finding of 50% of these being likely
low albedo classes, we estimate that 4% of the total discovered NEO population has low albedos based on these
factors alone. This is lower than the 9% value estimated by
Fernandez et al. (2001), although their sample considered
both NEOs and objects in unusual (but not near-Earth) orbits.
The observed population of NEOs having both T 3 and
measured (or inferred from taxonomic class) low albedos are
considered prime candidates for extinct comets, as reviewed
by Weissman et al. (1989, 2002). Deriving the actual population of objects having both T 3 and spectra presumably
compatible with extinct comet origins must take into account the discovery bias against these objects arising from
their inferred low albedos. Stuart (2003) (updated by Stuart
and Binzel, 2004) performs this diameter limited bias analysis and finds that 30 ± 4% of all NEOs have T 3, where
85 ± 15% of these have C-, D- or P-type taxonomies. These
two factors (0.30 × 0.85) yield an estimate of 25 ± 5% for
the total proportion of NEOs that have both T 3 and C-,
D-, and P-type spectroscopic properties above any given diameter. Limiting our attention to just D-types, for which the
diameter limited bias analysis of Stuart (2003) finds them
to be 43 ± 19% of all T 3 objects, yields an estimate of
13 ± 6% (equal to 0.30 × 0.43) for the total proportion of
NEOs that have both T 3 and D-type spectroscopic properties. We use this range of 13–25% for the percentage of
NEOs above any given size having both orbital and spectral characteristics being compatible with extinct comets as
a starting point for our discussion in Section 6.
5. Size dependence of spectral properties: evidence for
space weathering
While NEOs generally fall within regions of principal
component space (Fig. 1) populated by main-belt examples, many NEOs are uniquely found to populate the Qtype region and the intervening spectral component space
between the S- and Q-types (Binzel et al., 1996). Q-types
have long been associated as spectral analogs to ordinary
chondrite meteorites, with 1862 Apollo being the prototype example for this class (McFadden et al., 1985). The
possible link between the most common asteroids (S-types)
and the most common meteorites (ordinary chondrites) has
been long debated (e.g., Wetherill and Chapman, 1988;
Chapman, 1996). Thus, revealing a connection between the
S- and Q-type asteroids has the potential to establish the long
sought link between S-type asteroids and ordinary chondrite
meteorites.
To analyze this possible connection, we utilize the subset of near-Earth and Mars crossing objects having Q-, Sq-,
and S-types as determined using principal components measured within the SMASSII system (Bus, 1999) as compiled
here in Table 1. We restrict ourselves to consider the Qtypes, Sq-types, and the S-type core of the S-complex only.
The component showing the greatest difference between the
Q-, Sq-, and S-types is the Slope parameter, defined as the
average slope of the spectrum over the wavelength interval 0.44- to 0.92-µm (Bus, 1999). To explore a possible
size dependence, we convert the H magnitudes in Table 1
to diameters by utilizing the albedos determined through
the thermal modeling of this same population (Delbo et al.,
2003). For objects with no specific albedo determination, we
use the mean albedo (pv ) values (0.244 for S-types; 0.257
for Sq- and Q-types) from the Delbo et al. (2003) sample to
estimate diameters from:
D = 1329(pv )−1/2 10−0.2H .
While both random and systematic errors may be present
in the parameters used to estimate diameters, these errors
should have little effect in the broad statistical analysis we
apply here.
The data points in Fig. 7 display the diversity of spectral slopes versus diameters for Q-, Sq-, and S-types, where
Spectral properties of near-Earth objects
277
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-crossing
objects residing within the S-, Sq-, and Q-classes of Bus (1999). A running box mean is shown by filled squares (box size = 50, with error bars depicting
the 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-type
surfaces, depicted by the Q → S vector. (This vector corresponds to the Slope difference between the Q-type 1862 Apollo and the main-belt average for
S-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. The
SMFe model vectors were determined by calculating the slopes of meteorite spectra before and after applying the 0.05% SMFe curve from Pieters et al. (2000),
using the same method as Binzel et al. (2001b). All meteorite spectra are from Gaffey (1976) where “H” is the average for H6 chondrites, “L” is the Bald
Mountain L4 chondrite, and “LL” is the Olivenza LL5 chondrite.
the errors discussed above simply contribute to their scatter. Most evident is the higher dispersion for the smallest
objects, although we note that at larger sizes their fewer
numbers would tend to make any comparable dispersion less
apparent. To search for trends, we employ a running box
mean (box size 50) stepping through the sample one object at a time from smallest to largest. The resulting mean
values (with the overall sample variance reduced to the standard deviation of the mean) show a clear trend for decreasing
spectral slope with increasing size. While a size dependence
in spectral properties for S-types has been previously noted
(e.g., Gaffey et al., 1993a, 1993b; Rabinowitz, 1998), the
very small diameters sampled herein for the NEO population reveal a new characteristic: over the range of 0.1 to
5 km, mean spectral slopes appear to increase sharply and
then asymptotically approach a value of 0.44 near a diameter
of 5 km. Most interestingly, this value of 0.44 corresponds
to the mean spectral slope of all main-belt S-type asteroids
within the SMASSII sample of Bus (1999). One possible interpretation of Fig. 7 is that 5 km represents a “critical limit”
at which a size dependent transition for Q-type to S-type asteroids reaches “completeness.”
We first evaluate whether this trend is real or due to an
observational selection effect. Smaller objects must typically
be closer to the Earth in order to have their physical properties measured. Being closer, these objects can have a wider
range of phase angles (Earth–asteroid–Sun angles) than objects located farther away. An analysis to determine whether
phase angle induces any effect (such as reddening) on the
spectral slope was conducted by MIT undergraduate student Nancy Hsia. This analysis found no correlation within
our sample between phase angle and Slope. Similarly, no
correlation was found between Slope and the observed magnitude, thus revealing no systematic effect in our results as a
function of limiting magnitude for the observations.
We consider two possibilities for real effects that may
cause the trend in Fig. 7. The first is related to surface
particle sizes. Nominally, larger bodies should have finer
regolith characteristics owing to their greater gravity and
longer surface evolution lifetimes. (Larger bodies have a
longer lifetime in between collisions sufficiently energetic
to catastrophically disrupt them.) Thus we ask the question:
can the trend in Fig. 7 be due to decreasing particle size as
part of the process of regolith development? An analysis of
albedo and spectral contrasts within Psyche crater on Eros
by Clark et al. (2001) provides a case against regolith development causing the increasing spectral slope with increasing
diameter. For an olivine, orthopyroxene, plagioclase mixture intended to model Eros (an S-type asteroid), Clark et al.
(see their Fig. 18) found that finer grain sizes had decreased
spectral slopes relative to coarse grains. This model result is
opposite to the trend within Fig. 7.
The second possibility is that the age of the surface, as
affected by “space weathering” (see review by Clark et al.,
2002) causes the observed trend. Current models for space
weathering being due to the deposition of submicroscopic
iron (SMFe) (Hapke et al., 1975; Pieters et al., 2000) suggest that older surfaces become increasingly reddened with
278
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Fig. 8. Albedo versus diameter for S-, Sq-, and Q-class near-Earth and Mars-crossing objects. Data are from the tabulations by Binzel et al. (2002), Harris and
Lagerros (2002), and Delbo et al. (2003). Individual values are plotted by open circles with a running box mean based on n = 5 for the box size. (Error bars
depict 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).
the increasing accumulation of SMFe over time. Vectors
within Fig. 7 show the effect on the Slope parameter for H,
L, and LL chondrites by adding just 0.05% SMFe, where
the direction and magnitudes of these vectors are consistent
with the apparent transition from Q-type to S-type asteroids.
Small asteroids have shorter collisional lifetimes than larger
ones, suggesting that on average smaller asteroids will have
younger, fresher, and less reddened surfaces. Not all small
asteroids must have fresh surfaces since collisions are a stochastic process. However, the dispersion in surface properties may be expected to be higher at small sizes since they are
most likely to display “fresh” surfaces in addition to surfaces
that are stochastically “old.” Figure 7 indeed shows a greater
dispersion at the smaller sizes and a reddening trend (increasing spectral slope) with increasing size and presumably
older average surface age. If this trend is indeed an age effect, the most important implication for the space weathering
hypothesis is that the apparent limit at 5 km sets a timescale
over which the responsible process effectively reaches maturation. In other words, the average age of a 5 km asteroid
surface sets the time scale for the space weathering process
to be effective in “transforming” the spectral signatures of
ordinary chondrite like bodies into those of S-class asteroids.
Figure 8 suggests, but does not offer convincing evidence,
that 5 km is also a transition size for decreasing albedo with
increasing size. Revelation of an albedo trend, discussed by
Delbo 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 constraining
the extinct comet fraction, NEO source regions, and possible
space weathering trends.
Estimates for the percentage of extinct or dormant comets
within the NEO population have ranged from as high as
∼ 50% (Wetherill, 1988) to the currently predicted vicinity
of 2–10% (Bottke et al., 2002), where the latter estimate was
made for a specific magnitude range, 13 < H < 22. As described by Wiessman et al. (2002) (which incorporates the
Bottke et al. (2002) findings), physical studies combined
with dynamical parameters have provided key information
for constraining this percentage and for identifying individual objects as specific extinct comet candidates. Apart from
the goal of simply understanding the make-up of the NEO
population, identifying extinct comet candidates and determining their overall fraction provides an understanding of
the end states of comets.
Two quantities, Tisserand parameter and taxonomic class
(and/or measured albedo) are currently the best indicators
for recognizing bodies that may be extinct comets within
the NEO population. Current magnitude limited discovery
statistics show only 7% of all NEOs have T 3, and in Section 4 we found roughly one-half of these (and consequently
4% of the total observed NEO population) have taxonomic
properties consistent with low albedos. The combination of
eccentric orbits for T 3 objects and low albedos combine
as strong bias factors against their discovery and physical
characterization. Stuart (2003) (updated in Stuart and Binzel,
2004) performs a diameter limited bias correction analysis utilizing extensive NEO search statistics from LINEAR
(Stokes et al., 2002; Stuart, 2001), NEO albedos (Delbo et
al., 2003), and the taxonomic distributions presented here.
The resulting diameter limited estimate for the fraction of
the NEO population residing in T 3 orbits is 0.30 ± 0.04.
Diameter limited bias-corrected fractions for the T 3 population having low albedos within C-, D-, or P-type or just
D-type taxonomies are 0.85 ± 0.15 and 0.43 ± 0.19, respectively. All of these factors are summarized within columns
a–f of Table 4, where high and low values for these para-
0.43
0.62
0.24
0.85
1.00
0.70
0.30
0.34
0.26
Best estimate 0.07
High estimate 0.07
Low estimate 0.07
n
m
Extinct
comet
fraction C-,
D-, P-types
{g × j + k}
0.18
0.23
0.13
Correction
for extinct
comets
scattered to
T >3
+0.02
+0.04
0.00
k
j
Correction
for dark
asteroids
scattered to
T 3
0.65
0.75
0.55
Candidate
comet
fraction
D-types
{d × f}
0.13
0.19
0.07
h
g
Candidate
comet
fraction C-,
D-, P-types
{d × e}
0.25
0.31
0.19
f
T 3
fraction
having
D-types
e
T 3
fraction
having C-,
D-, P-types
Debiased
NEOs T 3
d
c
Inferred low
albedos
among
NEOs
{a × b}
0.04
0.05
0.03
Discovered
NEOs T 3
Description
of factors
b
a
Column label
SMASS
T 3
fraction
having C-,
D-, P-types
0.50
0.70
0.30
Final estimates
Dynamical factors
(Bottke et al., 2002 model)
Diameter limited bias corrected fractions
(Stuart, 2003, model)
Observed fractions
(magnitude limited)
Table 4
Observational and model parameters for estimating the extinct comet fraction within the NEO population
Extinct
comet
fraction
D-types
{h × j + k}
0.10
0.15
0.05
Spectral properties of near-Earth objects
279
meters are from the one-sigma uncertainties arising from
Poisson sampling statistics or from the formal error analyses
within the models. For all calculations progressing through
the columns of Table 4, these errors are propagated assuming they are independent and normally distributed.
There are a variety of ways to examine the factors within
Table 4 in order to achieve a new estimate for the extinct
comet fraction based on current discovery, albedo, and taxonomy statistics. Allowing that C-, D-, and P-type objects
are compatible in terms of spectra and inferred albedo with
extinct comet nuclei, then an estimated 25% (determined
from 0.30 × 0.85) of the NEO population above any given
diameter has properties consistent with being derived from
extinct comets. Restricting the extinct comet candidates to
those having the lowest inferred albedos (D-types), yields a
13% estimate (0.30 × 0.43) for candidate extinct comets in
the population. Factoring in the associated uncertainties (Table 4, columns g and h) expands the range for both estimates
considerably.
Having a T 3 orbit and taxonomic properties consistent with a low albedo, however, is not sufficient to imply an
object is an extinct comet. One of the complicating factors
is that low albedo asteroids can be scattered into T 3 orbits, thereby contaminating the sample of candidate extinct
comets. From the model of Morbidelli et al. (2002b) (also
A. Morbidelli, W. Bottke, personal communication, 2003),
for any given diameter or larger, 35% of the low albedo objects within the T 3 population are expected to be derived
from asteroids scattered mostly from the outer belt. Thus for
low albedo T 3 objects, the fraction being extinct comets
is estimated to be 0.65 ± 0.10. Correcting the T 3 candidate comet fraction (Table 4, columns g and h) by the 0.65
factor yields estimates of 8 to 16% for the diameter limited extinct comet fraction within the entire NEO population.
However, in reaching a final estimate, the fraction of the extinct comet population scattered from T 3 to T > 3 cannot
be ignored. Wiessman et al. (2002) (Table 1 and Table 3)
list many extinct comet candidates (such as 3200 Phaethon)
having Tisserand values greater than 3. We estimate the fraction of the extinct comet population scattered to T > 3 to be
0.02 ± 0.02, where this is an ad hoc value intended to recognize the potential existence of these scattered objects while
allowing that their numbers may be small. With this final
parameter we reach the two rightmost columns of Table 4
giving 10–18% as our best estimate for the NEO population, for any given diameter or larger, being comprised by
extinct comets. Propagating our uncertainties, formally this
estimate is 10 ± 5% if only D-type objects are considered as
comet candidates and 18 ± 5% considering all dark classes
(C, D, P).
How does our 10–18% range compare with the previous
estimates? Direct comparisons must be done carefully because of different debias criteria. Our results are diameter
limited, meaning above any given size, 10–18% of the NEO
population is estimated to be extinct comets. The 2–10% result of Bottke et al. (2002) is modeled over the magnitude
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R.P. Binzel et al. / Icarus 170 (2004) 259–294
range, 13 < H < 22. Accounting for the differences between
magnitude limited and diameter limited bias corrections (see
discussion by Morbidelli et al., 2002a) increases the range
given by Bottke et al. These authors (A. Morbidelli, W. Bottke, personal communication, 2003) find their diameter limited extinct comet contribution to be 17%, a value quite consistent with the estimate we report here.
While a compositional gradient has been well known
within the asteroid belt (Gradie and Tedesco, 1982) and a
variety of asteroidal source regions have been identified for
the NEO population (e.g., Wetherill and Chapman, 1988;
Greenberg and Nolan, 1989; Morbidelli et al., 2002a), their
chaotic routes to the inner Solar System (Wisdom, 1985)
may give little expectation of recognizable signatures linking NEOs to their sources. With the contributions of the
SMASS program and many other observers, the sample size
of spectral and taxonomic properties of the near-Earth and
Mars-crossing populations has now grown large enough to
be investigated for distinct signatures. Key to this investigation is the statistical model of Bottke et al. (2002) that
assigns probabilities for an object’s source based on its current orbital elements. The source region signatures revealed
within Fig. 6, with one exception (M-types, discussed below), show results consistent with the basics of the Gradie
and Tedesco (1982) compositional gradient within the asteroid belt. This consistency gives a mutual check between
the observational statistics from SMASS and other observers
and the source models of Bottke et al. (2002). E-types,
known to dominate the inner belt Hungaria region, show a
predominant inner belt source region signature, and a direct “tracer” may be most apparent for objects just across the
Mars-crossing boundary. (Table 3 lists candidates for additional objects that may be traced to distinct sources based on
orbital and spectral characteristics.) S-types show a clear signature from the inner main-belt where they are well known
to dominate. The indistinguishable source signature of the
Q-types with the S-types is a strong argument for their common origin, consistent with their spectral difference being an
artifact of a process such as space weathering (Clark et al.,
2002). C-types and P-types very nicely show their strongest
source region signatures from the outer asteroid belt, and the
D-type signature from Jupiter family comets constrains the
comet population, as discussed above.
The occurrence of V-types among NEOs but their rarity among Mars-crossers serves as a dynamical tracer
that specifically constrains the orbital eccentricity of their
source(s). While high eccentricity objects can evolve into
NEOs with a long (and likely to be observed) residence time
as Mars-crossers, objects in initially low eccentricity orbits
spend very little time as Mars-crossers in the course of their
“fast track” evolution through the ν6 and 3:1 resonances
(Morbidelli et al., 2002a). Thus low eccentricity sources appear the most likely source for the V-types, a condition met
by Vesta and its likely association of “Vestoids” (Binzel and
Xu, 1993). Based on the orbital distribution of V-types in
their own sample, Dandy et al. (2003) come to the same
“fast track” conclusion for the V-types.
As noted in Section 4, the M-types provide the most confounding source region signature in Fig. 6, suggesting a predominant origin from the outer belt. Such an origin does not
follow the colloquial “metallic” interpretation for “M,” more
likely to come from the more strongly heated region of the
inner asteroid belt. However, Jones et al. (1990) as well as
others (e.g., Rivkin et al., 1995, 2000, 2002) have found evidence of water of hydration in some M-types, implying any
“metallic” interpretation of their composition solely based
on visible wavelength data (and albedo) may be suspect.
(The mnemonic for “M” may be best characterized by the
word “muddle.”) While it is tempting to simply attribute this
outer belt spike for an M-type source to the revelation of “hydrated M-types” within the outer belt, we strongly caution
against any premature conclusion since this spike for the Mtypes within Fig. 6 is based on a sample of only three objects.
Strongly influencing this spike is the M-type object (6178)
1986 DA, whose orbital elements (a, e, i = 2.81, 0.58, 4.31)
lead to a Bottke et al. (2002) model probability of 0.81 for
coming from an outer belt source. 1986 DA is an object for
which the interpretation of high metal content seems secure,
as it exhibits one of the highest radar reflectivities measured
to date (Ostro et al., 1991). We suggest that the outer belt
source probability (if meaningful, since there is still a 19%
chance for coming from other regions) for the origin of 1986
DA may be a direct consequence of this object’s strength and
possible long-term dynamical evolution within the main-belt
prior to being cast into near-Earth space. Assuming 1986
DA is indeed “metallic,” it may have had a particularly long
main-belt lifetime against collisional disruption, as corroborated by the very long cosmic-ray exposure ages of iron
meteorites (Buchwald, 1975). As a counterpoint to the “most
traceable” objects being those just over the Mars-crosser
eccentricity border, the least traceable to their early Solar
System formation location may be high strength objects that
can have very long survival lifetimes and experience significant Yarkovsky drift within the main-belt prior to entering a
resonance transporting them to near-Earth space.
There is now a large body of spectral evidence linking
S-type asteroids to Q-type asteroids, implying a link between S-asterods and ordinary chondrite meteorites (e.g.,
Binzel et al., 1996; Chapman, 1996; Rabinowitz, 1998;
Trombka et al., 2000). First we caution that objects denoted
as “S-asteroids” span a wide range of mineralogies from
those being highly pyroxene-rich and olivine poor (the SI
sub-class of Gaffey et al., 1993a) to those being dominated
by olivine (the Gaffey SVII sub-class). For this reason, in
Section 5 we focused our search for S-type to Q-type relationships just along the transition in spectral slope from
Q-types to Sq-types to the core of the S-types so as to be
minimally effected by variations in mineralogy. We find a
size dependence to spectral properties, where Q-type asteroids are more common at smaller sizes, a result previously
noted by others, including Binzel et al. (1996), Hammergren
Spectral properties of near-Earth objects
(1998), Hicks et al. (1998), Rabinowitz (1998), Whiteley
(2001), and Dandy et al. (2003). Our conclusion that this
diameter dependent trend is largely caused by “space weathering” effects on spectral slope, rather than by variations in
mineralogy, stands to be verified or refuted by follow-up observations for these objects over near-infrared wavelengths.
At present, the consistency of the magnitude of the overall
change with the effects of a minor amount (0.05%) of submicroscopic iron (Fig. 7), adds weight to this particular model
for space weathering.
Our observations (Fig. 7) point to a key transition occurring at 5 km, where an increase in slope with increasing size
matches up with the average value for main-belt S-types.
(Main-belt S-types have been measured over the size range
of ∼ 10 to several hundred km.) An inflection in the running box trend occurring at 2 km may be the characteristic
in spectral colors noted by Rabinowitz (1998) based on an
analysis of filter photometry. At present it is not clear what
significance (if any) there is for the 2 km inflection as Fig. 7
shows that the dispersion of the raw data remains similar
over the full range from ∼ 100 m to 5 km. The transition at
5 km appears more robust in that there is a marked decrease
in dispersion of the raw data. The (perhaps now “classical”)
interpretation of a size dependent trend is that we are seeing
objects with increasing average collisional ages and therefore increasing average surface exposure times, i.e., increasingly “weathered” surfaces (Gaffey et al., 1993a; Binzel et
al., 1996, 1998; Rabinowitz, 1998). The observed transition
over the range of 0.1 to 5 km (Fig. 7) suggests we are seeing a “completion” of the space weathering process. If 5 km
represents the size at which the effects of space weathering become “complete,” then 5 km may represent a critical
size where the timescales for two independent processes are
matched. The first is the timescale over which collisions excavate and refresh the surface with unweathered material.
This timescale must be less than or equal to the collisional
disruption age for 5 km objects, for which estimates range
from 107 years (Farinella et al., 1998) to 109 years within
the main-belt (O’Brien and Greenberg, 2003), depending on
models for the body’s impact strength. Cheng (2004) argues
that objects larger than 5 km are survivors over the age of the
Solar System while those smaller are second (or later) generation fragments that must have younger collisional ages for
their surfaces. The second timescale is the interval for deposition of sufficient submicroscopic iron (as an example of a
space weathering process) to alter the slopes from ordinary
chondrite material to S-type asteroids.
The “classical” model of surface exposure time fails,
however, if the effective timescale for space weathering is
extremely short compared with collisional timescales. In
other words, if space weathering is so rapid as to be “instantaneous” compared with the interval over which an asteroid’s surface is refreshed, then surface exposure age is
irrelevant. (If the youngest surfaces of the smallest bodies
are instantly weathered, there is nothing about surface age
or collisional age that distinguishes them from the older
281
and also weathered surfaces of larger bodies.) Current models suggest that space weathering processes indeed may be
effective very rapidly, perhaps in as little as 50,000 years
(Hapke, 2001), a time that appears very short compared to
collisional timescales. We note that the possible very young
dynamical age of the Karin family (Nesvorny et al., 2002)
could provide insight into space weathering timescales.
For the case of rapid space weathering, the 0.1 to 5 km
trend in Fig. 7, with an apparent “completion” at 5 km, requires an alternate explanation from that of surface age. We
propose that the trend is a measure of increasing regolith development, driven both by the average surface age and the
increasing gravity (necessary for regolith retention) of the
body. (Generally decreasing rotation rate, with increasing
size, may also be a factor for regolith retention.) Under this
scenario, space weathering cannot commence until a regolith
begins to develop. The possible inflection at 2 km noted
above (Rabinowitz, 1998), could be the start of sufficient regolith development or retention and therefore the onset of
weathering effects. As the abundance of regolith grows, it is
instantly weathered, and creates surface reflectance properties that increasingly change from being ordinary chondritelike to being like S-asteroids. The “completeness plateau” at
5 km may be the size where there is sufficient surface evolution and gravity to begin to retain regolith (or a particular
particle size distribution for a regolith) in a manner that remains consistent with the regolith properties of 10–100 km
main-belt S-type asteroids.
7. Conclusion
The richness of our increasing scientific understanding of
the near-Earth object population is a natural product of the
necessity to characterize this population toward the practical goal of defining their size distribution and impact hazard
(Stuart, 2001, 2003; Stuart and Binzel, 2004). Our scientific
interest in the NEO population is driven by our desire to derive insights to their asteroidal and cometary sources and to
understand asteroid–meteorite connections. The NEO spectral data set is now growing large enough to correlate physical properties with dynamical source regions. The correlation of low albedo D-, C-, and P-types with Jupiter family
comet sources is an important synergy between dynamical
models (e.g., Bottke et al., 2002) and physical observations,
revealing a strong source signature for comets within the
NEO population. Taking into account the best available discovery, taxonomy, albedo, and bias correction models, we
estimate that 10–18% of the NEO population (at or above
any given diameter) may be extinct comets.
With the growing spectral data set and source region models, correlations between S-type asteroids and Q-type (ordinary chondrite-like) bodies continue to build. Both classes
have identical asteroid source region profiles. A clear sizedependent transition appears to forge a link between Sasteroids and ordinary chondrites, where surface exposure
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R.P. Binzel et al. / Icarus 170 (2004) 259–294
age, the process of regolith development, and/or regolith alteration (perhaps by submicroscopic Fe as the space weathering agent) completes the transition at sizes of 5 km. If
this interpretation is correct, a substantial proportion of Stype asteroids 5 km and above may be composed of ordinary chondrite-like materials, but have mature (reddened)
surfaces giving rise to spectral mismatches with ordinary
chondrites. In situ measurements by the NEAR-Shoemaker
spacecraft of the S-asteroid Eros (Trombka et al., 2000) give
the most direct and independent evidence for this link to ordinary chondrites.
Continued reconnaissance of the NEO population will reveal additional unusual compositions, raise new questions
for the detailed understanding of their origins, uncover numerous additional extinct comet candidates, and provide further direct evidence for the spectral evolution of asteroid
surfaces that is necessary to fully unravel remaining questions about asteroid–meteorite connections. Having a wellsampled population enables a scientific basis for choosing
future spacecraft mission targets among these relatively easily accessible worlds. Perhaps most pragmatically, all that
we learn scientifically is of direct practical benefit to the understanding of the long-term impact hazard to Earth.
Acknowledgments
We thank MIT students Lindsey Malcom, Nancy Hsia,
and April Deet Russell who were involved in various data
processing or early analysis stages. R.P.B. acknowledges
support for this research by NASA Grant NAG5-12355 and
NSF Grant AST-0205863 with additional funding support
from The Planetary Society. We thank many colleagues,
most especially W. Bottke and A. Morbidelli, for their many
helpful discussions that helped shape the ideas and results
presented here. We are grateful to A. Morbidelli and B.
Clark for their supportive and helpful reviews. We thank
the Bob Barr and the staff at the MDM Observatory where
this research originated. Binzel and Rivkin were Visiting
Astronomers at Kitt Peak National Observatory, National
Optical Astronomy Observatory, which is operated by the
Association of Universities for Research in Astronomy, Inc.
(AURA) under cooperative agreement with the National Science Foundation. Binzel, Burbine, and Stuart had the privilege to be Visiting Astronomers at the Infrared Telescope
Facility, which is operated by the University of Hawaii under
Cooperative Agreement no. NCC 5-538 with the National
Aeronautics and Space Administration, Office of Space Science, Planetary Astronomy Program. Observations obtained
at the Hale Telescope, Palomar Observatory are part of a
collaboration between the California Institute of Technology, NASA/JPL, and Cornell University. The work at the Jet
Propulsion Laboratory, Caltech, was supported under contract from NASA. All newly reported spectral data presented
here are available at http://smass.mit.edu/.
Appendix A
Observation summary
Number and name
433
433
433
433
433
433
719
1011
1036
1036
1620
1627
1627
1627
1862
1862
1864
1865
1866
1916
1917
1980
2062
2078
2078
2100
2100
2100
2102
2102
2201
2335
2340
2340
2423
3102
3102
3103
3103
3103
3103
3122
3122
3199
3199
3200
3288
3352
3352
3552
3552
3671
3671
3674
3691
3753
3753
3908
3908
3908
4034
4055
Eros
Eros
Eros
Eros
Eros
Eros
Albert
Laodamia
Ganymed
Ganymed
Geographos
Ivar
Ivar
Ivar
Apollo
Apollo
Daedalus
Cerberus
Sisyphus
Boreas
Cuyo
Tezcatlipoca
Aten
Nanking
Nanking
Ra-Shalom
Ra-Shalom
Ra-Shalom
Tantalus
Tantalus
Oljato
James
Hathor
Hathor
Ibarruri
Krok
Krok
Eger
Eger
Eger
Eger
Florence
Florence
Nefertiti
Nefertiti
Phaethon
Seleucus
McAuliffe
McAuliffe
Don Quixote
Don Quixote
Dionysus
Dionysus
Erbisbuhl
Bede
Cruithne
Cruithne
Nyx
Nyx
Nyx
Magellan
Provisional designation Observing date Telescope
1898 DQ
1898 DQ
1898 DQ
1898 DQ
1898 DQ
1898 DQ
1911 MT
1924 PK
1924 TD
1924 TD
1951 RA
1929 SH
1929 SH
1929 SH
1932 HA
1932 HA
1971 FA
1971 UA
1972 XA
1953 RA
1968 AA
1950 LA
1976 AA
1975 AD
1975 AD
1978 RA
1978 RA
1978 RA
1975 YA
1975 YA
1947 XC
1974 UB
1976 UA
1976 UA
1972 NC
1981 QA
1981 QA
1982 BB
1982 BB
1982 BB
1982 BB
1981 ET3
1981 ET3
1982 RA
1982 RA
1983 TB
1982 DV
1981 CW
1981 CW
1983 SA
1983 SA
1984 KD
1984 KD
1963 RH
1982 FT
1986 TO
1986 TO
1980 PA
1980 PA
1980 PA
1986 PA
1985 DO2
2-Sep-95
MDM 1.3m
3-Sep-95
MDM 1.3m
27-Oct-95
MDM 1.3m
2-Dec-95
IRTF 3m
9-Dec-95
MDM 2.4m
29-Jan-96
MDM 2.4m
23-Oct-01
Palomar 5m
20-Jan-02
KPNO 4m
5-Nov-94
MDM 1.3m
24-Feb-97
MDM 2.4m
7-Jan-94
MDM 2.4m
9-May-95
MDM 2.4m
10-Feb-97
IRTF 3m
25-Feb-97
MDM 2.4m
2-Dec-96
MDM 2.4m
4-Jan-97
IRTF 3m
20-Feb-94
MDM 2.4m
14-Oct-98
MDM 2.4m
6-Jan-94
MDM 2.4m
23-Oct-01
Palomar 5m
1-Apr-94
MDM 2.4m
28-May-97
MDM 2.4m
9-Feb-95
MDM 2.4m
2-Jan-97
IRTF 3m
10-Feb-97
IRTF 3m
4-Jan-97
IRTF 3m
13-Sep-97
MDM 2.4m
30-Sep-97
IRTF 3m
10-May-95
MDM 2.4m
12-May-95
MDM 2.4m
10-Dec-95
MDM 2.4m
3-Mar-00
IRTF 3m
23-Nov-97
MDM 2.4m
24-Nov-97
MDM 2.4m
3-Mar-00
IRTF 3m
6-May-00
IRTF 3m
6-Jul-00
Palomar 5m
31-Mar-94
MDM 2.4m
1-Apr-94
MDM 2.4m
10-Feb-97
IRTF 3m
25-Feb-97
MDM 2.4m
22-Jan-97
MDM 2.4m
10-Feb-97
IRTF 3m
10-Feb-97
IRTF 3m
24-Feb-97
MDM 2.4m
15-Nov-94
MDM 2.4m
24-Dec-01
Palomar 5m
22-Feb-94
MDM 2.4m
27-Feb-99
IRTF 3m
5-May-00
IRTF 3m
23-Oct-01
Palomar 5m
7-Apr-97
MDM 2.4m
25-May-97
MDM 2.4m
9-Aug-99
IRTF 3m
29-Jan-96
MDM 2.4m
13-Sep-97
MDM 2.4m
30-Sep-97
IRTF 3m
9-Sep-96
MDM 2.4m
12-Oct-96
MDM 2.4m
4-Jan-97
IRTF 3m
6-Aug-97
Keck 10m
1-Mar-00
KPNO 4m
(continued on next page)
Spectral properties of near-Earth objects
Appendix A (continued)
Number and name
4055
4055
4179
4183
4183
4197
4197
4341
4451
4503
4688
4947
4954
4954
4954
4954
4957
5131
5131
5143
5275
5587
5604
5626
5626
5641
5646
5660
5751
5751
5828
5836
5836
6047
6047
6047
6053
6455
6489
6489
6569
6611
7336
7336
7341
7358
7358
7480
7482
7822
7888
7889
7889
7977
8176
8566
9400
9400
10115
10165
10563
11311
11398
11398
Magellan
Magellan
Toutatis
Cuno
Cuno
Appendix A (continued)
Provisional designation Observing date Telescope
1985 DO2
1985 DO2
1989 AC
1959 LM
1959 LM
1982 TA
1982 TA
Poseidon
1987 KF
Grieve
1988 JJ
Cleobulus
1989 WM
1980 WF
Ninkasi
1988 TJ1
Eric
1990 SQ
Eric
1990 SQ
Eric
1990 SQ
Eric
1990 SQ
Brucemurray 1990 XJ
1990 BG
1990 BG
Heracles
1991 VL
Zdislava
1986 UU
1990 SB
1992 FE
1991 FE
1991 FE
McCleese
1990 DJ
1990 TR
1974 MA
Zao
1992 AC
Zao
1992 AC
1991 AM
1993 MF
1993 MF
1991 TB1
1991 TB1
1991 TB1
1993 BW3
1992 HE
Golevka
1991 JX
Golevka
1991 JX
1993 MO
1993 VW
Saunders
1989 RS1
Saunders
1989 RS1
1991 VK
Oze
1995 YA3
Oze
1995 YA3
Norwan
1994 PC
1994 PC1
1991 CS
1993 UC
1994 LX
1994 LX
1977 QQ5
1991 WA
1996 EN
1994 TW1
1994 TW1
1992 SK
1995 BL2
Izhdubar
1993 WD
Peleus
1993 XN2
1998 YP11
1998 YP11
283
3-Mar-00
4-Mar-00
20-Jan-97
23-Nov-97
8-Jan-98
11-Sep-96
12-Oct-96
25-May-97
20-Jan-02
25-Feb-99
16-Dec-00
12-Oct-96
24-Feb-94
28-Mar-94
8-Feb-97
2-Mar-00
1-Dec-96
7-Feb-95
27-Jan-99
12-Oct-96
10-Aug-99
17-May-01
1-Mar-00
19-Feb-94
1-Apr-94
2-Mar-00
28-Apr-96
22-Aug-93
17-Nov-94
9-Feb-95
22-Feb-02
21-Aug-93
1-Oct-97
17-Feb-98
15-Oct-98
23-Mar-99
9-Dec-95
8-Feb-95
9-May-95
22-May-99
10-May-95
7-Jan-94
11-Sep-96
12-Oct-96
12-Oct-96
29-Jan-96
15-Sep-98
1-Dec-96
19-Jan-97
24-Feb-97
28-Mar-94
2-May-98
16-May-01
8-Jan-98
11-Dec-95
28-Apr-96
14-Nov-94
7-Feb-95
27-Feb-99
9-Feb-95
11-Dec-95
5-Jan-94
24-Feb-99
27-Feb-99
IRTF 3m
IRTF 3m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
KPNO 4m
KPNO 4m
Palomar 5m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
IRTF 3m
MDM 2.4m
MDM 2.4m
IRTF 3m
MDM 2.4m
IRTF 3m
KPNO 4m
KPNO 4m
MDM 2.4m
MDM 2.4m
IRTF 3m
MDM 2.4m
MDM 1.3m
MDM 2.4m
MDM 2.4m
KPNO 4m
MDM 1.3m
IRTF 3m
IRTF 3m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
KPNO 4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
MDM 2.4m
MDM 2.4m
MDM 2.4m
KPNO 4m
IRTF 3m
Number and name
Provisional designation Observing date Telescope
11405
11405
11500
12538
12711
12711
12923
13651
13651
14402
14402
15745
15817 Lucianotesi
16657
16960
17274
17511
17511
18736
19356
20255
20255
20425
20790
20826
22099
22771
23548
23548
24475
25330
31345
31345
31346
32906
35107
35670
36284
37336
40310
48603
53319
1999 CV3
1999 CV3
1989 UR
1998 OH
1991 BB
1991 BB
1999 GK4
1997 BR
1997 BR
1991 DB
1991 DB
1991 PM5
1994 QC
1993 UB
1998 QS52
2000 LC16
1992 QN
1992 QN
1998 NU
1997 GH3
1998 FX2
1998 FX2
1998 VD35
2000 SE45
2000 UV13
2000 EX106
1999 CU3
1994 EF2
1994 EF2
2000 VN2
1999 KV4
1998 PG
1998 PG
1998 PB1
1994 RH
1991 VH
1998 SU27
2000 DM8
2001 RM
1999 KU4
1995 BC2
1999 JM8
1989 VA
1991 BN
1992 BF
1993 TQ2
1994 AB1
1994 AW1
1994 TF2
1995 WL8
1995 WL8
1996 BZ3
1996 FQ3
1996 FQ3
1997 AC11
1997 AQ18
1997 BQ
1997 BQ
1997 CZ5
1997 GL3
1997 RT
1997 RT
1997 SE5
25-Feb-99
KPNO 4m
27-Feb-99
IRTF 3m
15-Oct-98
MDM 2.4m
7-May-00
IRTF 3m
7-Feb-96
MDM 2.4m
15-Jan-00
IRTF 3m
23-May-99
IRTF 3m
8-Feb-97
IRTF 3m
25-Feb-97
MDM 2.4m
29-Feb-00
KPNO 4m
2-Mar-00
IRTF 3m
6-Jul-00
Palomar 5m
6-Aug-97
Keck 10m
7-Jan-94
MDM 2.4m
15-Oct-98
MDM 2.4m
6-Jul-00
Palomar 5m
15-Dec-95
MDM 2.4m
12-Oct-96
MDM 2.4m
17-May-01
KPNO 4m
10-Apr-97
MDM 2.4m
28-Mar-98
MDM 2.4m
30-Apr-98
IRTF 3m
17-Dec-00
Palomar 5m
16-Dec-00
Palomar 5m
16-Dec-00
Palomar 5m
5-May-00
IRTF 3m
23-Oct-01
Palomar 5m
30-Mar-94
MDM 2.4m
1-Apr-94
MDM 2.4m
16-Dec-00
Palomar 5m
16-Dec-00
Palomar 5m
18-Sep-98
IRTF 3m
13-Oct-98
MDM 2.4m
15-Sep-98
IRTF 3m
23-Feb-02
KPNO 4m
26-Feb-97
MDM 2.4m
6-Mar-02
Palomar 5m
22-Feb-02
KPNO 4m
27-Oct-01
KPNO 4m
24-May-99
IRTF 3m
9-Feb-95
MDM 2.4m
11-Aug-99
IRTF 3m
15-Nov-94
MDM 2.4m
28-Nov-02
Palomar 5m
8-Jan-98
MDM 2.4m
7-Jan-94
MDM 2.4m
20-Feb-94
MDM 2.4m
17-Dec-00
Palomar 5m
6-Aug-97
Keck 10m
14-Dec-95
MDM 2.4m
15-Dec-95
MDM 2.4m
9-Feb-96
MDM 2.4m
28-Apr-96
MDM 2.4m
30-Apr-96
MDM 2.4m
18-Jan-97
MDM 2.4m
18-Jan-97
MDM 2.4m
9-Feb-97
IRTF 3m
24-Feb-97
MDM 2.4m
9-Feb-97
IRTF 3m
10-Apr-97
MDM 2.4m
13-Sep-97
MDM 2.4m
30-Sep-97
IRTF 3m
23-Nov-97
MDM 2.4m
(continued on next page)
284
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Appendix A (continued)
Appendix A (continued)
Number and name Provisional designation Observing date Telescope
1997 TT25
1997 UH9
1997 US9
1998 FM5
1998 FM5
1998 HT31
1998 KU2
1998 KU2
1998 MQ
1998 QR15
1998 QR15
1998 SG2
1998 VO33
1998 VR
1998 WM
1998 WZ6
1998 XB
1998 XM4
1999 AQ10
1999 CV8
1999 CW8
1999 DB2
1999 DY2
1999 EE5
1999 FB
1999 FK21
1999 HF1
1999 KW4
1999 NC43
1999 NC43
2000 BG19
2000 BG19
2000 BG19
2000 BJ19
2000 CE59
2000 CK33
2000 CN33
2000 CO101
2000 CO101
2000 DO1
2000 DO1
2000 DO8
2000 EA107
2000 EZ148
2000 GD2
2000 GJ147
2000 GK137
2000 GO82
2000 GR146
2000 GU127
2000 GV127
2000 JG5
2000 JQ66
2000 KL33
2000 MU1
2000 NM
2000 OJ8
2000 PG3
2000 RW37
2000 SY162
2000 WC67
2000 WF6
2000 WJ10
2000 WJ63
23-Nov-97
23-Nov-97
23-Nov-97
28-Mar-98
30-Apr-98
2-May-98
16-Sep-98
13-Oct-98
9-Dec-98
18-Sep-98
13-Oct-98
15-Oct-98
8-Dec-98
9-Dec-98
8-Dec-98
8-Dec-98
27-Jan-99
17-Dec-00
27-Jan-99
24-Feb-99
24-Feb-99
24-Feb-99
24-Feb-99
23-Mar-99
23-Mar-99
22-Feb-02
23-May-99
23-May-99
1-Mar-00
3-Mar-00
1-Mar-00
4-Mar-00
6-May-00
17-Dec-00
1-Mar-00
21-Jan-02
1-Mar-00
1-Mar-00
2-Mar-00
1-Mar-00
3-Mar-00
3-Mar-00
5-Mar-02
5-May-00
5-Mar-02
5-May-00
6-Jul-00
7-May-00
5-May-00
6-May-00
5-May-00
7-May-00
6-Jul-00
6-Jul-00
6-Jul-00
6-Jul-00
17-Dec-00
20-Jun-01
7-Mar-01
17-Dec-00
18-Jan-01
17-Dec-00
17-Dec-00
16-Dec-00
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
IRTF 3m
IRTF 3m
MDM 2.4m
MDM 2.4m
IRTF 3m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
MDM 2.4m
IRTF 3m
Palomar 5m
IRTF 3m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
MDM 2.4m
MDM 2.4m
KPNO 4m
IRTF 3m
IRTF 3m
KPNO 4m
IRTF 3m
KPNO 4m
IRTF 3m
IRTF 3m
Palomar 5m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
IRTF 3m
KPNO 4m
IRTF 3m
IRTF 3m
KPNO 4m
IRTF 3m
KPNO 4m
IRTF 3m
Palomar 5m
IRTF 3m
IRTF 3m
IRTF 3m
IRTF 3m
IRTF 3m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
Magellan 6.5m
KPNO 4m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
Number and name
Provisional designation Observing date Telescope
2000 WK10
2000 WL10
2000 WL63
2000 WM63
2000 WO107
2000 XL44
2000 YA
2000 YH66
2000 YO29
2001 DU8
2001 EB
2001 EC
2001 FY
2001 HA8
2001 HK31
2001 HW15
2001 JM1
2001 JV1
2001 MF1
2001 OE84
2001 PD1
2001 QA143
2001 QQ142
2001 SJ262
2001 TC45
2001 TX16
2001 TY44
2001 UA5
2001 UC5
2001 UU92
2001 UY4
2001 VG5
2001 VS78
2001 WA25
2001 WG2
2001 WH2
2001 WL15
2001 XN254
2001 XR1
2001 XS1
2001 XS30
2001 XU30
2001 XY10
2001 YE1
2001 YK4
2002 AA
2002 AD9
2002 AH29
2002 AK14
2002 AL14
2002 AQ2
2002 AU5
2002 AV
2002 BA1
2002 BK25
2002 BM26
2002 BP26
2002 CS11
2002 CT46
2002 DH2
2002 DO3
2002 DY3
2002 EA
2002 EC
6-Mar-02
16-Dec-00
16-Dec-00
17-Dec-00
16-Dec-00
7-Mar-01
17-Dec-00
18-Jan-01
18-Jan-01
7-Mar-01
17-May-01
7-Mar-01
16-May-01
16-May-01
16-May-01
16-May-01
16-May-01
17-May-01
23-Dec-01
23-Oct-01
28-Oct-01
23-Oct-01
6-Mar-02
27-Oct-01
27-Oct-01
27-Oct-01
24-Dec-01
23-Feb-02
23-Oct-01
24-Dec-01
27-Oct-01
24-Dec-01
21-Jan-02
24-Dec-01
24-Dec-01
23-Dec-01
24-Dec-01
22-Feb-02
20-Jan-02
24-Dec-01
23-Dec-01
24-Dec-01
24-Dec-01
24-Dec-01
20-Jan-02
21-Jan-02
21-Jan-02
21-Jan-02
20-Jan-02
6-Mar-02
5-Mar-02
20-Jan-02
20-Jan-02
6-Mar-02
22-Feb-02
22-Feb-02
23-Feb-02
5-Mar-02
22-Feb-02
5-Mar-02
6-Mar-02
5-Mar-02
6-Mar-02
5-Mar-02
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
KPNO 4m
Palomar 5m
Palomar 5m
Palomar 5m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
Palomar 5m
Palomar 5m
KPNO 4m
Palomar 5m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
Palomar 5m
KPNO 4m
Palomar 5m
Palomar 5m
KPNO 4m
Palomar 5m
KPNO 4m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
KPNO 4m
KPNO 4m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
Palomar 5m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
Palomar 5m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
KPNO 4m
Palomar 5m
KPNO 4m
Palomar 5m
KPNO 4m
Spectral properties of near-Earth objects
285
Appendix B
Spectra for near-Earth and Mars-crossing asteroids. These data are available in digital format at http://smass.mit.edu/
(continued on next page)
286
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Appendix B (continued)
(continued on next page)
Spectral properties of near-Earth objects
287
Appendix B (continued)
(continued on next page)
288
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Appendix B (continued)
(continued on next page)
Spectral properties of near-Earth objects
289
Appendix B (continued)
(continued on next page)
290
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Appendix B (continued)
(continued on next page)
Spectral properties of near-Earth objects
291
Appendix B (continued)
(continued on next page)
292
Appendix B (continued)
R.P. Binzel et al. / Icarus 170 (2004) 259–294
Spectral properties of near-Earth objects
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