Cometary ices were altered in the early solar nebula ?
Abundance Medians and Lower and Upper Quartile Values of Ices with Respect to Water ice
(Oberg et al. 2011)81
CO/CO
2 ratioinsufficient CO2 sublimation ?
(Oberg et al. 2011)
protostellar samples
related to outburst events ?
(Ootsubo et al. 2012) 82
Summary for comet ice
AKARI/IRCによる彗星18天体の近赤外線 (2.5‒5 μm) 観測 AKARI/IRC は H2O, CO2, CO を同時に観測した
これまでで最大の彗星 CO2 観測データベース CO2
18彗星のうち 17 天体で検出 (29P/SW1 at 6 AU 以外) CO2/H2O は < 4--30 % の幅でばらついている
長周期・短周期で明確な差は見られていない（サンプル少だが）
CO 検出はわずかに 3 彗星 (29P, C/2006 W3, C/2008 Q3) CO/H2O 比はほとんど上限値のみ得られている
彗星における CO2/H2O, CO/CO2
- CO2 : 低質量星周囲の氷よりも depleted
- 微惑星形成時点での熱的要因（彗星核の形成場所を示唆？）
- CO/CO2 < 1 --- 原始太陽系星雲は酸化的か？
83
( Morbidelli & Levison 2003 )
長周期彗星と 短周期彗星
・
氷もダストも内部の組成の違いは大きくない・それぞれのカテゴリー内でのばらつきが大きい
・それほど遠い領域で形成されたわけではなさそう and/or
・原始太陽系星雲中で物質が十分にかき混ぜられていた？
短周期彗星
長周期彗星
黄道光・黄道光放射
• Interplanetary dust (particles) (惑星間塵)
• 惑星間空間に広く分布する sub-μm〜mmサイズの 固体微粒子
• 特に中間-遠赤外線波長帯では非常に卓越した前景光
• 彗星、小惑星起源、外縁天体起源、および星間空間起源
• 太陽系小天体及び塵の最近の力学的進化を知る手がかり
• 系外惑星系・系外黄道光との比較対象という点でも重要
DIRBE mission-averaged sky maps (in Ecliptic coordinates) (Kelsall+ 1998)
惑星間塵の粒径分布
• Interplanetary dust (particles) (惑星間塵)
• 惑星間空間に広く分布する sub-μm〜mmサイズの
固体微粒子 FLUX OF INTERPLANETARY DUST 249
2 i i i f i i i i ~
0 " ~ - - - inte'planetary flux-
- 2 ~ - ~ ' ~ i ° n ' e r 8 9 ... I . . . . flux
~ -6 ~
-10 -12 -14 -16
-18 I I I I I I I [ I
-18 -16 -14 -12 -10 - 8 - 6 ~ - 2 0 2
log m (g)
FIG. 1. Cumulative particle flux on a spining flat plate at 1 A U distance from the Sun. T h e plate normal vector lies in a n d the spin axis is perpendicular to the ecliptic plane. For an isotropic flux the stated flux val- ues for a flat plate c o r r e s p o n d to an effective solid angle of ~r sr. T w o flux models (interplanetary and lunar flux) a n d the flux of fl meteoroids which results from our analysis of collisional effects are shown. T h e lunar flux curve has been c o n s t r u c t e d from the relative lunar crater size distribution and the absolute fluxes from the P e g a s u s spacecrafts. T h e interplanetary flux curve deviated from the lunar flux for small particles, taking into a c c o u n t the H E O S 2 and the fl-meteoroid fluxes, both m e a s u r e d by Pioneers 8 and 9 and calcu- lated.
ber of techniques which include photo- graphic and radar meteor data, capacitor or pressurized cell penetration sensors, and space exposed impact sensors--such as spacecraft windows. More recently, impact ionization detectors for the measurement of small (m < l0 9 g) particles have been suc- cessfully flown. An important synthesis and summary of many of the earlier, but still valid, data and their intercorrelation is given by Naumann (1966). Here we adopt his fit II for the Pegasus penetration data from the 0.04- and 0.02-cm-thick detectors which we consider highly reliable and well analyzed. These data are representative also of other large meteoroid (m - l 0 - 7 g )
spacecraft data which include Explorer 16 and 23 results and optical and radar meteor data. More recently, e.g., Cour-Palais (1974), supported Naumann's results.
Another set of meteoroid flux data exists for small meteoroids (m -< 10 -9 ), that is, the flux information and anisotropy infor- mation of the interplanetary dust flux from
the HEOS-2 (Hoffmann et al., 1975a, b) and from the Pioneer 8 and 9 (Berg and Griin, 1973) dust experiments. The detailed method of how we arrive at each of the two flux models is described in the appendix.
c. F l u x C u r v e s
Figure 1 shows the relative size distribu- tion of the steepest lunar microcrater data (lunar flux) as well as the interplanetary flux. Both fluxes are required to fit the sat- ellite data at large meteoroid masses. It is evident that both sets of satellite data (Peg- asus and HEOS-2) cannot be made to per- fectly fit the lunar relative size distribution.
There is hardly an overlap within the esti- mated errors of the measurement.
We have adjusted the meteoroid flux curve through the large particle (Pegasus) fluxes for the following reason: Flavill et al.
(1978) and Allison and McDonnell (1981) examined the effects of secondary micro- craters produced by ejecta from primary craters on lunar samples. They concluded that secondary microcratering has a signifi- cant effect in the 1- to 10-/zm-diameter cra- ter range. The magnitude of the secondary cratering effect depends on the impact ge- ometry. Based on the experimental data available to them (Schneider, 1975; Flavill and McDonnell, 1977), they estimate the number of secondary to primary craters to be of the order of unity in the relevant size range (1- to 10-tzm-diameter crater). Zook
et al. (1984) reported results from new hy- pervelocity impact experiments which showed that the number of secondary im- pact pits is more than 2 orders of magni- tude higher than was previously thought (Schneider, 1975). The basic difference be- tween the new experiments and the one by Schneider is that the new experiments used oblique impact angles in contrast to normal impacts. Since oblique impacts are more re- alistic for lunar impacts we have to accept the conclusion by Zook et al. (1984) that
"the lunar impact pit population, for pit di- ameters below about 7 micrometers, is probably dominated by high-speed second-
256 GROGAN, DERMOTT, AND DURDA
FIG.8. The iterative filtering procedure. Panel (a) shows a raw model dust band having the same viewing geometry as an observed background (b). In the first iteration (a) is added to (b) and the sum is filtered to obtain (c), a model dust band (smooth curve); the observed dust bands (noisy curve) are also shown for comparison. The background obtained from this iteration shown in panel (d) is of a higher intensity than the original background because it contains two low-frequency dust band components, one from the addition of the model dust band and one from the actual dust band in the original observed background. In the final iteration, we subtract the excess intensity shown in (d) from the original background (b) and add (a) before filtering to obtain the final dust band (e) and the final background (f) that agree with the observations.
and filter a second time, the new low-frequency background should now have a higher amplitude, reflecting the fact that it now contains two low-frequency dust band components, one from the observations and one from the model. The difference in amplitude between the two backgrounds therefore gives us an estimate of the extent of the contribution of the dust bands to the low-frequency zodiacal background. In essence, by using the same filter in the modeling process that we use to define the observed dust bands and by iterating, we are able to bypass the arbitrary divide associated with the filter, and extract the un- derlying low-frequency component of the dust bands that other techniques are unable to retrieve. This is essential in revealing the true extent to which asteroidal dust contributes to the cloud.
4. THE NATUREOFTHE SIZE–FREQUENCY DISTRIBUTION OFASTEROIDALDUST
This work differs from our previous modeling of the dust bands (Grogan et al. 1997) in that our models include a size–
frequency distribution, rather than being composed of particles of a single size. This is critical in our efforts to provide a model of the dust bands that can match the IRAS observations in multiple wavebands. Particles ranging in size from 1 to 100 µm are in-
cluded, each of which are assumed to be Mie spheres composed of astronomical silicate (Draine and Lee 1984). The lower end cut-off is determined by the fact that contribution to the thermal emission from particles smaller than this size is negligible. The upper cut-off follows from the fact that in the zodiacal cloud, the P–R drag lifetime is comparable to the collisional lifetime for particle of diameter 100–200 µm (Zook and McKay 1986, Leinert and Gr¨un 1990, Nishiizumi et al. 1991, Flynn 1992).
We realize that particles even larger than this will exist in the zodiacal cloud, but we have not obtained the dynamical history for these particles. We continue to work on the dynamics of par- ticles up to and beyond 500 µm, but in this regime we will have to start incorporating the effects of particle–particle collisions as the P–R drag timescales become longer than the collisional lifetimes. This is a topic for a future paper.
There is a wealth of evidence for the existence of large par- ticles in the zodiacal cloud. Gr¨un et al. (1985) review such evi- dence from a variety of sources including lunar microcratering and spacecraft micrometeroid detectors (Helios, Pioneer), com- bined with assumed dust particle scattering functions. The con- clusion is that the bulk of zodiacal emission is produced from particles in the order of tens to hundreds of microns in size.
Nishiizumi et al. (1991) measure the space exposure times of large particles retrieved from the Greenland ice cap and conclude that the results are consistent with P–R drag lifetimes from the main-belt, and have radiogenic isotopic ratios inconsistent with a high eccentricity (assumed cometary) origin. There is also the evidence from the LDEF cratering record (Love and Brownlee 1993), shown in Fig. 9, which suggests a peak in the particle
FIG. 9. The terrestrial influx of zodiacal dust particles, as measured from the cratering record on the LDEF satellite. The slope of area against particle mass indicates a value forq, the size–frequency distribution index, of approximately 1.2.
(Grogan+2001)LDEF
(Gruen+85)
惑星間塵の供給源：彗星
•
1 AU付近では、おそらく短周期彗 星がメインの供給源(Nesvorny+2011)
•
彗星ダストトレイルの発見など、大 きめ（> 100μｍ）のダストが惑星 間空間に供給されていることが分 かってきたEncke by ISO
Reach et al. (2000)
Sarugaku et al. (2007)
惑星間塵の供給源：小惑星
• 小惑星ダストバンド
(Low et al. 1984, Sykes et al. 1988)•
メインベルト内の小惑星族での衝突起源？•
ここ最近（< 107年）程度以内の衝突•
1.4 band•
Beagle family origin ? (<15 Myr ago)•
2.1 band•
Karin cluster origin ? (5.8 Myr ago)•
10 ( 9.35 ) band•
Veritas family origin ? (8.3 Myr ago)(Nesvorny et al. 2003)
leading trailing
No data
No data
52 S Y K E S A N D G R E E N B E R G
q
(0)
0 °
o : a ( . 0
180 ° Ecliptic Longitude
(b)
0 Ecliptic Lotitude
(c)
360 °
t
J
Fic;. 1. (a) The appearance of the zodiacal dust bands can be modeled by considering lhe superposition of particle orbits having essentially identical orbital elements with the exception of the longitude of the ascending node. In the top frame, the orbits of an ensemble of particles have similar orbital elements describe a sine wave in ecliptic coordinates. Each particle spends more lime at higher ecliptic latitudes than near the ecliptic. As the longitudes of nodes begin to disperse (middle) this tendency results in the beginning of band formation at maximum and minimum latitudes, initially as two dashes on either side of the ecliptic, whose centers are 180"
apart in ecliptic longitude. Finally, when the longitudes of nodes are lotally randomized (bollom), two parallel bands extend completely around lhe ecliplic. Excepl for the case of parlicle orbils having zero inclination, zodiacal dust bands would always be fi)und in pairs. (b) The variation in surface brighlness as a function of ecliptic latitude is found to have two sharp peaks al ecliptic lalitudes corresponding to Ihe mean orbital inclina- lion of the constituent particles. The sharpness of the peaks depends upon the dispersion of orbital elements among band pair particles. The greater the dispersion, lhe thicker the bands. (c) Debris from a single collisional event in the asteroid bell undergoes orbital evolution to form band pairs. Initially, debris escapes into similar heliocentric orbits. Variations in orbital velocities resull in their complete distribution in mean orbital phase in 100 to 1000 years (lop). Gravitational perturbations by Jupiter (whose orbil is shown as the thick ring) act on the collisional debris to precess their orbits. Variations in semimajor axes of debris orbits result in different preces- sion rates and an increasing dispersion in the longitudes of nodes over the ensemble of debris orbits (middle).
After 10 ~ to 10 ~ years, the nodes have been distributed completely around the ecliptic, forming a complete zodiacal band pair.
n a r y d e t e r m i n a t i o n s o f t h e s o l a r d i s t a n c e s o f t h e b a n d p a i r s s h o w n o e v i d e n t c o r r e l a - t i o n w i t h t h e m e a n d i s t a n c e s o f t h e a s t e r o i d f a m i l i e s . B r i g h t n e s s t e m p e r a t u r e s o f t h e
o u t e r b a n d p a i r , d e t e r m i n e d a t a s i n g l e e c l i p t i c l o n g i t u d e , y i e l d s o l a r d i s t a n c e s o f 3.2 A U f o r t h e n o r t h e r n b a n d a n d 2.2 A U f o r t h e s o u t h e r n b a n d , w h i l e t h e i n n e r b a n d
Sykes+Greenberg (1986)
18 & 90 μm)band
ZL observed with AKARI
Leading direction
Trailing direction
Pyo et al. 2010 Ootsubo et al. 2015
leading
• clear detection of ±1.4° and ±10° dust bands in far-IR
• trailing direction <10% brighter than leading direction
trailing
Observed by AKARI L18W-band
(18μm)
Leading direction
Trailing direction
α, β γ
γ C
J M
N
α, β γ C
J
M
N γ
D M
D D
Pyo et al. 2009