惑星大気(2)

地球惑星圏物理学 

惑星大気(2)

1

大気散逸

2

大気の宇宙空間への流出

大気量・組成 , 表層水量に影響を及ぼす

熱的散逸： Hydrodynamic escape, Jeans escape

非熱的散逸：光化学的散逸 , ion pickup, sputtering など

熱的散逸

Jeans escape

Hydrodynamic escape

Exobaseからの高速度粒子の“蒸発”

大気の流体的な“流出”

Exobase

3

半値幅 

(kT/m)

^1/2

4

非熱的散逸

Carr and Head (2003) estimated the volume of a potential early martian water reservoir from a geomorphological anal- ysis of possible shorelines. The data suggest that the surface features can be explained by a primitive (post-Noachian) martian water ocean equivalent to a global layer with a depth of about 150 m (Carr and Head, 2003).

The absence of valleys in the younger areas of the north- ern hemisphere suggests that the liquid water presumably present during the Noachian epoch disappeared suddenly, before or during the end of the Late Heavy Bombardment about 3.8 Ga. Based on the calculation by Kass (2001) that an equivalent global water layer with a depth of about 50 m may have escaped to space since the Hesperian epoch (about 3.5 Ga), and assuming that an amount of water equivalent to a global ocean with a depth of about 20–30 m may still be retained in the present polar caps, Carr and Head (2003) es- timated a subsurface water reservoir to be equivalent to a global ocean with a depth of about 80 m, which may be trapped in volatile-rich deposits on the surface or in a groundwater system.

It is important to note that after the Noachian epoch the martian atmosphere was most likely not protected by a strong intrinsic magnetic field. Therefore, the solar wind plasma could interact directly with the upper atmosphere.

Later studies by Leblanc and Johnson (2002) and Lammer et al. (2003a) of the water loss since the Hesperian epoch, after the martian magnetic dynamo had ceased operating, indi- cated that the earlier water loss estimates of 30–80 m (e.g., Luhmann et al., 1992; Zhang et al., 1993a; Kass and Yung, 1995, 1996, 1999; Kass, 2001; Krasnopolsky and Feldman, 2001) may be overestimations. The water loss study by Lam- mer et al. (2003a) included thermal escape, ion pickup, dis- sociative recombination, and sputtering, and used data of solar proxies with different ages to reconstruct the Sun’s ra- diation and particle environment from the present time back to 3.5 Ga (Dorren and Guinan, 1994; Guinan and Ribas, 2002;

Wood et al., 2002; Ribas et al., 2005). They estimated a total loss of water equivalent to a global martian ocean with a depth of about 12 m over the past 3.5 billion years by as- suming the self-regulation mechanism between the loss of O and H, as postulated by McElroy and Donahue (1972).

However, as pointed out by Lammer et al. (2003a), some additional complex ionospheric processes, such as detached plasma clouds and cold ion outflow (Fig. 1), might also have contributed to the erosion of the atmosphere. The ejection of plasma clouds triggered by plasma instabilities or other pro- cesses at the ionopause transition region, which were ob- served at Venus (Brace et al., 1982) and studied by a number of authors (Wolff et al., 1980; Terada et al., 2002, 2004; Ar- shukova et al., 2004), should be considered an additional way by which heavy ions may have escaped from the martian ionosphere (Penz et al., 2004). Cold ion outflow from the ionosphere into the martian ionospheric plasma tail, due to momentum transfer from the solar wind plasma flow (e.g., Pérez-de-Tejada, 1992, 1998; Lundin et al., 1989, 1990, 1991, 2007, 2008a, 2008b; Terada et al., 2002, 2004), is another po- tentially important process. It would be logical to suggest that the cold ion outflow into the plasma tail could have been a very efficient loss process on Mars, because the same loss process was recently identified as one of the most efficient loss processes from the Venus atmosphere by ASPERA-4 ob- servations on board Venus Express (Barabash et al., 2007a).

Pérez-de-Tejada (1992) estimated that cold ion escape due to momentum transport effects may have removed an amount of water from Mars that would have been equiva- lent to a global ocean with a depth of about 10–30 m, de- pending on uncertainties in the solar wind parameters and the martian plasma environment in the past. Lammer et al.

(2003b) found, for moderate solar activity, a cold O^!ion out- flow loss rate from present-day Mars on the order of "10²⁵ s^#1, which is in agreement with the recently observed AS- PERA-3 ion escape rates into the plasma tail (Lundin et al., 2007, 2008a, 2008b). Assuming D/H isotopic constraints over the last 3.5 billion years, Lammer et al. (2003b) estimated a total water loss from Mars equivalent to a global ocean with a depth of about 14 m (minimum) to 34 m (maximum). To derive this estimate, they added together the loss rates due to O^!ion pickup (Lammer et al., 2003a), dissociative recom- bination loss of hot O* (Luhmann, 1997), O^!ion loss induced by plasma instabilities (Penz et al., 2004), and sputtering of O, CO, and CO2(Leblanc and Johnson, 2002). The main un- certainties in this estimate result from present uncertainties in the atmospheric and ionospheric parameters and the re- lated atmospheric escape rates due to plasma instabilities and cold ion outflow in the transition layer between the so- lar wind and the martian ionosphere. It should be noted that in previous estimations of cold ion outflow loss, Pérez-de- Tejada (1992) and Lammer et al. (2003b) assumed that the ions escaped through the entire 100% circular ring area around the terminator. Therefore, the estimated cold ion out- flow loss rates by these authors should be considered to be an upper limit.

An important question regarding the loss of the martian atmosphere and its initial water inventory involves the on- TERADA ET AL.

56

FIG. 1. Illustration of ion pickup, cold ion outflow, plasma clouds, and sputtering regions around Mars. The dashed line corresponds to the planetary obstacle (ionopause) and the solid line to the bow shock. !i is the thickness of the ionos- phere at the terminator.

太陽風と惑星大気

Terada et al. (2009)

を改変

Photochemical escape

窒素の例

9030 Fox AND DALGARNO.' NITROGEN ESCAPE FROM MARS

TABLE 2. Vibrational Distribution of N 2 + at 200 km

Relative Population

v" No O Quenching With O Quenching

0 57% 60%

1 21% 20%

2 10% 10%

3 5.6% 5.2%

4 3.4% 3.1%

5 2.0% 1.9%

excited and since dissociative recombination is faster for vi-

brationally excited ions, the isotope effect discussed by Wallis [1978] operates only 55% of the time. Quenching by col- lisions with atomic oxygen has a negligible effect on the vibra-

tional distribution at the exobase.

3.2. Escape Fluxes

Table 3 summarizes the escape fluxes we obtain for both low and high solar activity and their average. Dissociative

recombination of N2 + produces the largest escape flux of about 2.9 x 105 cm -2 s -•. Photodissociation produces an escape flux of about 1 x 105 cm-2 s-x, reaction (2) 9 x 10 '•

cm-2 s-1 and the other sources less than 4 x 10 '• cm-2 s-x.

Reaction (7) gives rise to an escape flux of less than 1 x 10 '•

cm-2 s-•. The total average escape flux is 5.6 x 105 cm -2 s-x. The values depend only weakly on the choice of electron

temperatures. For an electron temperature profile which ex- ceeds the ion temperatures and which reaches 2000 K near 250 km, the total average escape flux is enhanced by less than

15%.

4. VARIATION OF THE ESCAPE FLUXES WITH TIME

The escape fluxes varied in different ways as the atmosphere evolved. Leovy [1982] has concluded that the assumption of a constant homopause density level on Mars is probably valid.

If so, both the turbopause and exobase were higher at earlier times when the atmosphere contained more nitrogen.

Figure 3 shows the escape fluxes as a function of the total mixing ratio of N 2 relative to CO 2 for an arbitrary escape flux at the current N 2 number density. The escape fluxes due to photodissociation, to dissociative ionization of N2, and to dissociative recombination of N2 + are proportional to the concentration of N 2 at the exobase and for small mixing ratios of N 2 relative to CO 2 to the total column density of N 2. In the limiting case when the fraction of nitrogen in the atmosphere approaches unity, the column density of N2 will determine the altitude of the exobase and the escape flux will become con- stant. The escape fluxes due to chemical sources are pro-

TABLE 3. •'•N Escape Fluxes

Process

Low Solar

Flux

High

Solar

Flux Average

N 2 + hv--• N + N

N 2 + hv--• N + + N + e N 2 q- e---• N + + N + 2e N2 + + e--•N + N

N 2 + + O-• NO + + N

N 2 + O + + --} N + + N + O +

Total

5.9 x 10 '•

1.7 x 10 '•

1.4 x 10 '•

8.4 x 10 '•

4.7 x 10 '•

4.8 x 103 2.3 x 105

1.4 x 105 6.0 x 10 '•

4.8 x 10 '•

5.0 x 105 1.3 x 105 1.4 x 10 '•

8.9 x 105

1.0 x 105 3.0 x 10 '• 3.1 x 10 '• 2.9 x 105 8.9 x 10 '• 9.4 x 103 5.6 x 105 In cm -2 s -1

7.0

6.0-

5.0-

,4.0-

:5.0-

2.0-

1.0-

_I0 ^-2

[ [ ] I Ill] I I I I I I Ill I I I [ I Ill

½,, _

_

i• I io ø io I

[ N2] / [CO:> ]

Fig. 3. Relative escape fluxes as a function of the N2/CO 2 mixing ratio at the turbopause. The flux at the current value of the N 2 mixing ratio, 0.026, is arbitrarily set equal to 1.0; • is the flux pro- portional to the N 2 density at the exobase, •2 is the flux due to the reaction N 2 + + O--• NO + + N, and •b3 is the flux due to the reaction N 2 + O++•N + + N + O +.

portional to O or O + + as well as to N 2. The rate of reaction

(2) will achieve a maximum as the concentration of N 2 is increased but then begins to decrease as N 2 becomes the dominant species at the exobase and O decreases. The reac-

tion of O + + with N 2 is still more severely damped. As N 2 increases, loss of O + + becomes controlled by N 2. Because the O + + density is then inversely proportional to N2, the rate of

reaction (7) becomes proportional to the atomic oxygen den- sity. Thus the escape fluxes due to chemical reactions are rela- tively less important at earlier times.

5. INITIAL COLUMN DENSITIES AND ISOTOPE RATIOS

The rate of change of the column abundance, N, of N 2 is

dN/dt = --ok(t)

where &(t) is the escape flux. The equation for the time evolu-

tion of the ratio for XSNX'•N to X'•NX'•N is

df/dt = &(t)f(1 - R)N

where R is the net relative escape rate [McElroy et al., 1976;

Nier et al., 1976]. R depends on the ratio •SN•'•N/•'•N•'•N at

the exobase, which is about 0.82 [McElroy et al., 1977], and on the fraction of the escape flux due to dissociative recombi-

nation of ground state N 2 + ions.

By numerical integration over the lifetime of the planet,

4.5 x 109 years, we obtain an initial column density of 1.3 x 1023 cm -2. The calculated isotope enhancement ratio is

2.51, much larger after measured value of 1.62. It appears then that the escape rate was smaller at some time in the past than we predict on the basis of the current value.

Agreement with the measured isotope ratio can be obtained by postulating arbitrarily that the initial outgassing occurred

only about 1-2 x 109 years ago. This late outgassing may not

be necessary, however. The isotope enhancement ratio of 2.51

was calculaied for an average value of recently measured solar

fluxes whereas the intensity may have been different in the

分子反応を介してエネルギーを得る

Ion pickup

太陽風磁場が大気イオンを絡めとる

Sputtering

Pickupイオンが中性粒子を叩き出す

Cold ion flow

太陽風と電離大気の 

Kelvin-Helmholtz不安定

Fox & Dalgarno (1983)

Hydrodynamic escape

5

流体力学的 ( 静水圧平衡ではない ) 大気散逸 上層大気が加熱され、大気存在条件

を満たさない状況で発生 太陽風の惑星アナログ

地球型惑星の円盤ガス捕獲大気 , 金星の水 , 短周期系外惑星

48 第 4 章 惑星大気

(4.7) より、この大気保持の条件式は次式で与えられる。

λ 0 > γ

γ − 1 . (4.10)

式 (4.10) は、ある大気組成を与えた場合、惑星が大気を保持するために満たすべき地表面温

度の上限値 T esc を与える式と考えることもできる。表 4-1 に太陽系内の惑星・衛星の大気の 有無、及び T esc と平衡温度 T eff の比較がある。式 (4.10) は大気が存在する必要条件になって いるが、十分条件ではないことがわかる。

表 4-1 ．惑星・衛星の大気の有無。朝倉書店『惑星の科学』 1 章 ( 阿部豊著 ) より。

※

ここでの

λ

0は上層大気における

escape parameter

Image credit: NASA/CXC/M. Weiss

Escaping atmosphere

Transit light curve of exoplanet GJ 436b (Ehrenreich et al. 2015)

Some close-in exoplanets show UV transit depth deeper than visible transit (only 5 samples so far, because UV transit observations require space telescope)

Their absorption cross sections are larger than Hill radius It suggests that atmosphere is escaping!

8 Figure 2 | Lyman- α transit light curves of GJ 436b. a, b, Data are from visit 1 (circles), visit 2 (stars), visit 3 (squares) and visit 0 (triangles). All uncertainties are 1 σ . a, The Lyman- α (Ly α ) line is integrated over

[−120,−40] km s

⁻¹

and shows mean absorption signals with respect to the out-of-transit flux of 17.6 ± 5.2% (pre- transit), 56.2 ± 3.6% (in-transit) and 47.2 ± 4.1% (post-transit). b, The line is integrated over [+30,+200] km s

⁻¹

and shows no significant absorption signals: 0.7 ± 3.6% (pre-transit), 1.7 ± 3.5% (in-transit) and 8.0 ± 3.1%

(post-transit). With a depth of 0.69%, the optical transit (thin black lines in a and b) is barely seen at this scale between its contact points (dotted lines in a and b). A synthetic light curve (green) calculated from the three- dimensional numerical simulation

²⁰

is overplotted on the data in a.

UV(Lyman-α)

~ 10 %

V is ib le ~ 1 %

Time →

10 Figure 4 | Polar view of three-dimensional simulation representing a slice of the comet-like cloud coplanar with the line of sight. Hydrogen atom velocity and direction in the rest frame of the star are represented by

arrows. Particles are colour-coded as a function of their projected velocities on the line of sight (the dashed

vertical line). Inset, zoom out of this image to the full spatial extent of the exospheric cloud (in blue). The planet orbit is shown to scale with the green ellipse and the star is represented with the yellow circle.

Star Planet Comet-like tail?

6

恒星活動の時間進化

7

62

_第

4

_{章 惑星大気}

4.5 大気散逸

惑星大気の宇宙空間への流出過程を大気散逸と呼ぶ。大気散逸過程は惑星の歴史を通じて 大気や表層水の量、大気組成の進化に影響を及ぼしてきたと考えられている。大気散逸は局所 熱力学平衡にある気体が流出する熱的散逸・それ以外の非熱的散逸に大別することができる。

4.5.1 Hydrodynamic escape

Hydrodynamic escape

は熱的散逸に分類される大気散逸過程であり、上層大気が宇宙空間

に流体力学に流出する過程である。これは基本的には太陽風の惑星版と理解することもでき る。前節で述べたように、上層大気は恒星の短波長放射を吸収して高温となるため、下層大 気は大気を保持する条件式

(4.10)

を満たしていながら、上層大気はこれを満たさないという 状態が存在しうる。このような状況では、上層大気は静水圧平衡状態を維持することができ ず、上層大気は流体力学的に宇宙空間に流出していく。

Hydrodynamic escape

による大気流 出率を求めるためには上層大気の温度を決定する複雑なエネルギー収支を計算する必要があ るが、その上限値はエネルギー供給率、すなわち上層大気への恒星極紫外線フラックスによっ て制約される

(

_{エネルギー律速}

)

。また、下層大気は二酸化炭素など重い分子で構成されてい ながら、上層大気は水素など軽い分子で構成されており、

hydrodynamic escape

_{が駆動され} ているような状況も存在しうる。この場合、下層大気から上層大気への水素の輸送率が流出 率を制限することもある

(

_拡散律速

)

_。

現在の太陽系の惑星・衛星においては

hydrodynamic escape

は発生していないが、金星 の高い水素

/

_重水素

(D/H)

比など間接的な証拠から、過去には

hydrodynamic escape

_が発生 していたと考えられている。その物理的な主因は、太陽を含む恒星の極紫外線光度は若い時 代ほど大きいことである

(

_図

4-13)

。強力な極紫外線放射に晒された上層大気は高温となり、

hydrodynamic escape

が駆動される。また、形成後数十億年が経過した惑星系においても、恒

星のごく近傍に惑星が存在する場合、その惑星からは

hydrodynamic escape

_{による大気散逸} が発生しうる。実際に、軌道長半径が

0.1 AU

程度以下の太陽系外惑星について、宇宙望遠鏡 を用いた紫外線トランジット観測から

hydrodynamic escape

を示唆する観測結果が得られて いる。

図

4-13

．太陽アナログ星の観測に基づく極紫外線フラックスの時間進化。

Ribas et al. (2005)

より。

the predictions that equations (4) and (6) make for the mass-loss rates of very young stars should be considered questionable.

A theoretical point that should be made regarding equa- tions (1), (2), and (3) is that while the connections between V rot , t, and F X in equations (2) and (3) are rather direct, the physical connection between M M _ and F X suggested by equa- tion (1) must be indirect because the X-ray flux presumably originates in closed field regions while the stellar wind arises in open field regions. The apparent correlation between M M _ and F X in Figure 7 may mean that the underlying cause of stronger closed field regions that yield higher F X (i.e., the dynamo) also yields stronger open fields that accelerate more wind.

Figure 8 shows what relation (4) implies for the mass-loss history of the Sun, once again assuming that the mass loss

‘‘ saturates ’’ when the coronal X-ray flux saturates. The solar wind may have been ! 1000 times stronger when the Sun was very young, although we note once again that pre- dictions for very early times are suspect. The upper limits in Figure 8 are based on nondetections of radio emission from three solar-like stars (Gaidos, Gu¨del, & Blake 2000), illus- trating that the high mass-loss rates predicted for young stars are still consistent with our inability to detect radio emission from these stars. Note that the dependence of mass loss on time derived empirically in relation (4) likely results from a number of physical processes, including the evolu- tion of magnetic fields in both strength and structure (i.e., the relative importance of closed versus open structure).

Equation (4) has important ramifications for our under- standing of the angular momentum evolution of cool stars because it is the interaction between the stellar wind mate- rial and the magnetic field of the rotating star that is believed to provide the braking mechanism for stellar rota- tion on the main sequence, which is ultimately why rotation and activity both decrease with time. Higher mass-loss rates can be expected to yield stronger braking. Models for how magnetic braking should affect stellar rotation suggest rela- tions of the form

! _

!

! / M M _ M

R A R

! " m

; ð 7 Þ

where ! is the angular rotation rate and R A is the Alfve´n radius (Weber & Davis 1967; Stepien´ 1988; Gaidos, Gu¨del,

& Blake 2000). The exponent m is a number between 0 and 2, where m ¼ 2 corresponds to a purely radial magnetic field. Mestel (1984) argues that more reasonable magnetic geometries suggest m ¼ 0 1.

We are interested in the time dependence of the quantities in equation (7). Equation (2) shows that ! / V rot / t % 0:6 , so

! _

! / t % 1:6 and ! !=! _ / t % 1 . We assume that neither the stellar mass (M ) nor the radius (R) are time variable (see x 3.3), while equation (4) indicates the time dependence of M M _ . The Alfve´n radius is

R A ¼

ffiffiffiffiffiffiffiffiffiffiffiffi V w M M _

B 2 r s

; ð 8 Þ

where V w is the stellar wind speed and B r is the radial mag- netic field. Note that since equation (7) is computed assum- ing a simple, global magnetic field geometry, B r should be considered a disk-averaged field when one relates it to the far more complex fields that stars like the Sun actually have.

We assume that V w does not vary, which was also an assumption used in the derivation of mass-loss rates from the astrospheric absorption (see x 2.2). We can then express B r as a power law of time, B r / t ! , where from equations (4), (7), and (8) we find

! ¼ 1

m % ð 1:00 & 0:26 Þ ð m þ 2 Þ

m : ð 9 Þ

Assuming that m is in the physically allowable range of m ¼ 0 2 yields the upper limit ! < % 1:0, while the more likely range of m ¼ 0 1 suggested by Mestel (1984) implies

! < % 1:2. In any case, this result suggests that disk-aver- aged stellar magnetic fields decrease at least inversely with age for solar-like stars.

3.3. Relevance for Planetary Atmospheres

Figure 9 shows the cumulative mass loss for the Sun as a function of time based on the relation in Figure 8. Despite the high mass-loss rates predicted for the young Sun, the total mass that has been lost by the solar wind is still d 0.03 M ( , not enough to have dramatically altered the solar lumi-

Fig. 8.—Mass-loss history of the Sun suggested by the power-law rela- tion from Fig. 7. The upper limits are based on radio nondetections of three solar-like stars (Gaidos et al. 2000).

Fig. 9.—Cumulative mass loss for the Sun as a function of time, based on the mass-loss history of the solar wind in Fig. 8.

No. 1, 2002 MEASURED MASS-LOSS RATES OF SOLAR-LIKE STARS 423

短波長放射 (Ribas et al. 2005) 恒星風 (Wood et al. 2002)

若い恒星ほど紫外線放射・恒星風フラックスが大きい

→ 若い惑星ほど大気散逸率が大きい

  ↔ 可視・赤外放射フラックス：時間的に増加

  (40 億年前の太陽光度は現在の約 70%)

Jeans escape

8

静水圧平衡にある大気の外気圏界面 (exobase) からの流出

多くの場合、流出率は小さい

4.5. 大気散逸 63

4.5.2 Jeans escape

Jeans escape は上層大気まで静水圧平衡が成り立っている惑星大気からの熱的散逸過程で

ある。上層大気が全体として重力的に束縛されている場合でも、ある確率で一部の気体分子は 脱出速度を超える熱運動速度を持つ。下層大気においては気体分子の平均自由行程が短く、こ のような気体分子が散逸することはないが、十分上層においては、他の気体分子と衝突せず宇 宙空間に飛び去ってしまう。粒子が無衝突で大気圏外へ出てしまう領域を外気圏 (exosphere) と呼び、その下端エクソベース (exobase) の位置 r e は次式で定義される。

! ∞

r e

σ n(r) = 1. (4.18)

ここで σ は粒子の衝突断面積、 n は粒子数密度である。

次に Jeans escape による大気流出率を見積もる。気体分子の速度分布はマクスウェル分布

に従うとして、速さ v から v + dv の間にある粒子の密度は、

n(r e )f (v )dv = 4π n(r e )

" m g 2π k B T

# 3/2

exp

"

− m g v 2 2k B T

#

v 2 dv. (4.19)

エクソベースにおいて、単位時間・単位面積あたりに上向きに通過する粒子フラックスは、

1

4 n(r e )f (v )vdv = π n(r e )

" m g 2πk B T

# 3/2

exp

"

− m g v 2 2k B T

#

v 3 dv, (4.20) と書くことができる。ここで、ファクター 1/4 はエクソベースを下向きに通過する寄与がない こと、及び面に対して粒子の運動方向のなす角 θ の補正項をつけて立体角積分した寄与によ るものである。実際に宇宙空間へ飛び去ることができる粒子は脱出速度を上回る速さをもっ た粒子のみなので、 v > v esc について積分することで Jeans escape の粒子フラックスを得る。

F e =

" g 2π r e

# 1/2 1

σ (λ e + 1)λ 1/2 e exp( − λ e ). (4.21) ここで、 λ e は エクソベースでのエスケープ・パラメータである。また、 n(r e ) = σk m g g

B T (r e ) を 用いた。 具体的に地球の場合においてこの式を用いて粒子フラックスを計算すると λ e が 40 以上で

は 1 mol/yr 以下である。一般に、 Jeans escape による大気散逸量は大気質量と比べると十分

小さく、惑星が大規模に大気を失う過程は前述の hydrodynamic escape である。

単位面積・単位時間あたりの流出率

添字 e は exobase を表す

阿部豊

,

朝倉書店『惑星の科学』

1

章

, 1993

64

第

4

章 惑星大気

4.5.3

_{非熱的散逸}

図

4-14.

非磁化惑星における大気散逸の模式図。

Terada (2014), J. Plasma Fusion Res. Vol.90, No.12, 75

‐

79.

局所熱力学平衡にある大気の熱エネルギーによる散逸である熱的散逸に対し、これ以外の 散逸過程を総称して非熱的散逸と呼ぶ。非熱的散逸には外気圏における光化学反応による散 逸や太陽風によるイオンピックアップ、スパッタリングなどの多数の過程が存在する。太陽 風に起因する過程については、惑星磁場の有無が散逸率に大きく影響しうることがその特徴 である。非熱的散逸過程においては

Jeans escape

ではほぼ散逸しない重い元素も流出する可 能性があり、特にかつて火星に存在した

(

と考えられている

)

厚い大気の散逸過程として重要 視されている。

非熱的散逸

9

光化学反応 , 太陽風との相互作用

熱的散逸では殆ど失われない重い元素の散逸 火星大気

64 第4章 惑星大気

4.5.3

_{非熱的散逸}

図4-14. 非磁化惑星における大気散逸の模式図。Terada (2014), J. Plasma Fusion Res. Vol.90, No.12, 75‐79.

局所熱力学平衡にある大気の熱エネルギーによる散逸である熱的散逸に対し、これ以外の 散逸過程を総称して非熱的散逸と呼ぶ。非熱的散逸には外気圏における光化学反応による散 逸や太陽風によるイオンピックアップ、スパッタリングなどの多数の過程が存在する。太陽 風に起因する過程については、惑星磁場の有無が散逸率に大きく影響しうることがその特徴 である。非熱的散逸過程においてはJeans escapeではほぼ散逸しない重い元素も流出する可 能性があり、特にかつて火星に存在した(と考えられている)厚い大気の散逸過程として重要 視されている。

まとめ

10

大気散逸：惑星大気への宇宙空間の流出減少 熱的散逸： Hydrodynamic escape, Jeans escape

非熱的散逸：光科学的散逸 , Ion pick-up, sputtering など

若い恒星の紫外線放射・恒星風大 → 惑星の大気散逸率大 例：地球型惑星の円盤ガス捕獲大気 ( 一次大気 ),

   金星・火星の水 , 火星の初期大気