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Water: from clouds to planets
水:从云到行星

Ewine F. van Dishoeck  Ewine F. van Dishoeck Leiden Observatory, Leiden University, The Netherlands; Max Planck Institute for Extraterrestrial Physics, Garching, Germany
莱顿天文台,莱顿大学,荷兰; 马克斯·普朗克天体物理研究所,加尔兴,德国
   Edwin A. Bergin  Edwin A. Bergin University of Michigan, USA
密歇根大学,美国
   Dariusz C. Lis  Dariusz C. Lis California Institute of Technology, USA
美国加州理工学院
   Jonathan I. Lunine  乔纳森·I·卢宁 Cornell University, USA  美国康奈尔大学
Abstract  摘要

Results from recent space missions, in particular Spitzer and Herschel, have lead to significant progress in our understanding of the formation and transport of water from clouds to disks, planetesimals, and planets. In this review, we provide the underpinnings for the basic molecular physics and chemistry of water and outline these advances in the context of water formation in space, its transport to a forming disk, its evolution in the disk, and finally the delivery to forming terrestrial worlds and accretion by gas giants. Throughout, we pay close attention to the disposition of water as vapor or solid and whether it might be subject to processing at any stage. The context of the water in the solar system and the isotopic ratios (D/H) in various bodies are discussed as grounding data point for this evolution. Additional advances include growing knowledge of the composition of atmospheres of extra-solar gas giants, which may be influenced by the variable phases of water in the protoplanetary disk. Further, the architecture of extra-solar systems leaves strong hints of dynamical interactions, which are important for the delivery of water and subsequent evolution of planetary systems. We conclude with an exploration of water on Earth and note that all of the processes and key parameters identified here should also hold for exoplanetary systems.
最近太空任务的结果,特别是斯皮策和赫歇尔,已经显著推动了我们对水从云到盘、行星体和行星的形成和运输的理解取得了重大进展。在这篇综述中,我们提供了水的基本分子物理和化学的基础,并概述了这些进展,涉及了太空中水的形成背景、其运输到形成盘的过程、在盘中的演化,最终交付给形成地球类行星并被气态巨行星吸积。在整个过程中,我们密切关注水作为蒸汽或固体的状态,以及在任何阶段是否可能受到处理的问题。太阳系中水的背景和各种天体中的同位素比(D/H)作为这一演化的基础数据点进行讨论。其他进展包括对外星气态巨行星大气成分的日益增长的了解,这可能受到原行星盘中水的不同相影响。 此外,外星系结构留下了动力学相互作用的强烈暗示,这对水的输送和行星系统的后续演化至关重要。我们最后探讨了地球上的水,并指出这里确定的所有过程和关键参数也应适用于系外行星系统。

 
 
 

1 INTRODUCTION  1 引言

With nearly 1000 exoplanets discovered to date and statistics indicating that every star hosts at least one planet (Batalha et al., 2013), the next step in our search for life elsewhere in the universe is to characterize these planets. The presence of water on a planet is universally accepted as essential for its potential habitability. Water in gaseous form acts as a coolant that allows interstellar gas clouds to collapse to form stars, whereas water ice facilitates the sticking of small dust particles that eventually must grow to planetesimals and planets. The development of life requires liquid water and even the most primitive cellular life on Earth consists primarily of water. Water assists many chemical reactions leading to complexity by acting as an effective solvent. It shapes the geology and climate on rocky planets, and is a major or primary constituent of the solid bodies of the outer solar system.
迄今为止已发现近 1000 颗系外行星,并且统计数据表明每颗恒星至少拥有一颗行星(Batalha 等人),我们在宇宙中寻找其他生命的下一步是对这些行星进行表征。行星上存在水被普遍认为是其潜在适居性所必需的。气态水作为冷却剂,使星际气云坍缩形成恒星,而冰水促进小尘埃颗粒的粘附,最终形成小行星和行星。生命的发展需要液态水,即使地球上最原始的细胞生命主要由水组成。水通过作为有效溶剂参与许多化学反应,促进复杂性的产生。它塑造了岩石行星的地质和气候,并且是外太阳系固体天体的主要或主要组成部分。

How common are planets that contain water, and how does the water content depend on the planet’s formation history and other properties of the star-planet system? Thanks to a number of recent space missions, culminating with the Herschel Space Observatory, an enormous step forward has been made in our understanding of where water is formed in space, what its abundance is in various physical environments, and how it is transported from collapsing clouds to forming planetary systems. At the same time, new results are emerging on the water content of bodies in our own solar system and in the atmospheres of known exoplanets. This review attempts to synthesize the results from these different fields by summarizing our current understanding of the water trail from clouds to planets.
行星中含水的普遍程度如何,水含量如何取决于行星的形成历史和恒星-行星系统的其他属性?多亏了一系列最近的太空任务,最终以赫歇尔空间天文台为高潮,我们对水在空间中形成的位置有了更深入的了解,以及在各种物理环境中的丰度,以及它是如何从坍缩的云层运输到形成行星系统的。与此同时,有关我们太阳系中天体和已知外行星大气中水含量的新结果正在涌现。这篇综述试图通过总结我们对从云到行星的水迹的当前理解,综合这些不同领域的结果。

Speculations about the presence of water on Mars and other planets in our solar system date back many centuries. Water is firmly detected as gas in the atmospheres of all planets including Mercury and as ice on the surfaces of the terrestrial planets, the Moon, several moons of giant planets, asteroids, comets and Kuiper Belt Objects (see review by Encrenaz, 2008). Evidence for past liquid water on Mars has been strengthened by recent data from the Curiosity rover (Williams et al., 2013). Water has also been detected in spectra of the Sun (Wallace et al., 1995) and those of other cool stars. In interstellar space, gaseous water was detected more than 40 years ago in the Orion nebula through its masing transition at 22 GHz (1 cm; Cheung et al., 1969) and water ice was discovered a few years later through its infrared bands toward protostars (Gillett and Forrest, 1973). Water vapor and ice have now been observed in many star- and planet-forming regions throughout the galaxy (reviews by Cernicharo and Crovisier, 2005; Boogert et al., 2008; Melnick, 2009; Bergin and van Dishoeck, 2012) and even in external galaxies out to high redshifts (e.g., Shimonishi et al., 2010; Lis et al., 2011; Weiß et al., 2013). Water is indeed ubiquitous throughout the universe.
关于火星和太阳系其他行星上存在水的猜测可以追溯到许多世纪以前。水被确定为所有行星的大气中的气体,包括水星,并且在地球类行星、月球、几颗巨行星的卫星、小行星、彗星和柯伊伯带天体的表面上被检测到为冰(见文献综述)。最近来自好奇号火星车的数据加强了有关火星过去液态水存在的证据。太阳的光谱中也检测到了水,以及其他冷星的光谱。在星际空间中,超过 40 年前在猎户座星云中通过其 22 GHz(1 厘米)的微波跃迁检测到了水的气态形式,几年后通过其红外波段朝原恒星的方向发现了水冰。如今在银河系中已经观测到了许多形成恒星和行星的区域中的水蒸气和冰(文献综述),甚至在外部星系中也观测到了高红移的水(例如)。水确实是宇宙中无处不在的。

On their journey from clouds to cores, the water molecules encounter a wide range of conditions, with temperatures ranging from 10 K in cold prestellar cores to <<\sim2000 K in shocks and the inner regions of protoplanetary disks. Densities vary from 104\sim 10^{4} cm-3 in molecular clouds to 101310^{13} cm-3 in the midplanes of disks and 101910^{19} cm-3 in planetary atmospheres. The chemistry naturally responds to these changing conditions. A major question addressed here is to what extent the water molecules produced in interstellar clouds are preserved all the way to exoplanetary atmospheres, or whether water is produced in situ in planet-forming regions. Understanding how, where and when water forms is critical for answering the question whether water-containing planets are common.
从云层到行星核心的旅程中,水分子会遇到各种条件,温度从寒冷的前恒星核心的 10K 到冲击波和原行星盘的内部区域的 2000K 不等。密度从分子云中的 cm 到盘面中的 cm 和行星大气中的 cm 不等。化学反应自然地响应这些变化的条件。这里要解决的一个主要问题是,星际云中产生的水分子在通往外行星大气的过程中是否会被保留,或者水是否是在行星形成区域内产生的。了解水是如何、在哪里以及何时形成的对于回答含水行星是否普遍存在的问题至关重要。

2 H2O PHYSICS AND CHEMISTRY
2H 2 O 物理和化学

This section reviews the basic physical and chemical properties of water in its various forms, as relevant for interstellar and planetary system conditions. More details, examples and links to databases can be found in the recent review by van Dishoeck et al. (2013).
本节回顾了水在其各种形式中的基本物理和化学性质,这些性质与星际和行星系统条件相关。有关更多详细信息、示例和数据库链接,请参阅 van Dishoeck 等人(2013 年)最近的综述。

2.1 Water phases  2.1 水的相态

Water can exist as a gas (vapor or ‘steam’), as a solid (ice), or as a liquid. At the low pressures of interstellar space, only water vapor and ice occur, with the temperature at which the transition occurs depending on density. At typical cloud densities of 10410^{4} particles cm-3, water sublimates around 100 K (Fraser et al., 2001), but at densities of 101310^{13} cm-3, corresponding to the midplanes of protoplanetary disks, the sublimation temperature increases to \sim160 K. According to the phase diagram of water, liquid water can exist above the triple point at 273 K and 6.12 mbar (1017\sim 10^{17} cm-3). Such pressures and temperatures are typically achieved at the surfaces of bodies of the size of Mars or larger and at distances between 0.7 and 1.7 AU for a solar-type star.
水可以存在为气体(蒸汽或“蒸汽”)、固体(冰)或液体。在星际空间的低压下,只有水蒸气和冰存在,转变发生的温度取决于密度。在典型云密度为 104superscript10410^{4} 粒子 cm -3 时,水升华温度约为 100K(Fraser 等人,2001 年),但在密度为 1013superscript101310^{13} cm -3 ,对应于原行星盘的中平面时,升华温度增加到 similar-to\sim 160K。根据水的相图,液态水可以存在于三相点以上的 273K 和 6.12mbar( 1017similar-toabsentsuperscript1017\sim 10^{17} cm -3 )。这样的压力和温度通常在火星大小或更大的天体表面以及太阳型恒星 0.7 到 1.7 AU 之间的距离上实现。

Water ice can take many different crystalline and amorphous forms depending on temperature and pressure. At interstellar densities, crystallization of an initially amorphous ice to the cubic configuration, IcI_{c}, occurs around 90 K. This phase change is irreversible: even when the ice is cooled down again, the crystal structure remains and it therefore provides a record of the temperature history of the ice. Below 90 K, interstellar ice is mostly in a compact high-density amorphous (HDA) phase, which does not naturally occur on planetary surfaces (Jenniskens and Blake, 1994). The densities of water ice in the HDA, LDA and IcI_{c} phases are 1.17, 0.94 and 0.92 gr cm-3, respectively, much lower than those of rocks (3.2–4.4 gr cm-3 for magnesium-iron silicates).
水冰可以根据温度和压力采取许多不同的结晶和非晶形式。在星际密度下,初始非晶冰结晶为立方体结构, IcsubscriptI_{c} ,大约在 90K 左右发生。这种相变是不可逆的:即使冰再次冷却,晶体结构仍然保持,因此它提供了冰的温度历史记录。在 90K 以下,星际冰主要处于紧凑的高密度非晶(HDA)相,这在行星表面上并不自然发生(Jenniskens 和 Blake,1994)。HDA、LDA 和 IcsubscriptI_{c} 相中水冰的密度分别为 1.17、0.94 和 0.92 克/立方厘米,远低于岩石的密度(镁铁硅酸盐的密度为 3.2-4.4 克/立方厘米)。

Clathrate hydrates are crystalline water-based solids in which small non-polar molecules can be trapped inside ‘cages’ of the hydrogen-bonded water molecules. They can be formed when a gas of water mixed with other species condenses111Strictly speaking, the term condensation refers to the gas to liquid transition; we adopt here the astronomical parlance where it is also used to denote the gas-to-solid transition. out at high pressure and has enough entropy to form a stable clathrate structure (Lunine and Stevenson, 1985; Mousis et al., 2010). Clathrate hydrates are found in large quantities on Earth, with methane clathrates on the deep ocean floor and in permafrost as the best known examples. They have been postulated to occur in large quantities on other planets and icy solar system bodies.
克拉沙水合物是结晶的基于水的固体,其中小的非极性分子可以被困在氢键结合的水分子的“笼子”内。当水气与其他物种混合在高压下凝结并具有足够的熵形成稳定的克拉沙结构时,它们可以形成。克拉沙水合物在地球上大量存在,其中深海底部的甲烷水合物和冻土层中的甲烷水合物是最为人熟知的例子。据推测,它们在其他行星和冰冷的太阳系天体上也以大量存在。

2.2 Water spectroscopy  2.2 水光谱学

Except for in-situ mass spectroscopy in planetary and cometary atmospheres, all information about interstellar and solar system water comes from spectroscopic data obtained with telescopes. Because of the high abundance of water in the Earth’s atmosphere, the bulk of the data comes from space observatories. Like any molecule, water has electronic, vibrational and rotational energy levels. Dipole-allowed transitions between electronic states occur at ultraviolet (UV) wavelengths, between vibrational states at near- to mid-infrared (IR) wavelengths, and between rotational states from mid- to far-IR and submillimeter wavelengths.
除了行星和彗星大气中的原位质谱外,有关星际和太阳系水的所有信息都来自使用望远镜获得的光谱数据。由于地球大气中水的丰富,大部分数据来自空间天文台。像任何分子一样,水具有电子、振动和转动能级。电偶极允许的电子态之间的跃迁发生在紫外(UV)波长处,振动态之间的跃迁发生在近红外到中红外(IR)波长处,转动态之间的跃迁发生在中红外到远红外和亚毫米波长处。

Interstellar water vapor observations target mostly the pure rotational transitions. H2O is an asymmetric rotor with a highly irregular set of energy levels, characterized by quantum numbers JKAKCJ_{K_{A}K_{C}}. Because water is a light molecule, the spacing of its rotational energy levels is much larger than that of heavy rotors, such as CO or CS, and the corresponding wavelengths much shorter (0.5 mm vs 3–7 mm for the lowest transitions). The nuclear spins of the two hydrogen atoms can be either parallel or anti-parallel, and this results in a grouping of the H2O energy levels into ortho (KA+KC=K_{A}+K_{C}=odd) and para (KA+KC=K_{A}+K_{C}=even) ladders, with a statistical weight ratio of 3:1, respectively. Radiative transitions between these two ladders are forbidden to high order, and only chemical reactions in which an H atom of water is exchanged with an H-atom of a reactant can transform ortho- to para-H2O and vice versa.
星际水蒸气观测主要针对纯旋转跃迁。H2O 是一个不对称转子,具有高度不规则的能级集,其特征是量子数。因为水是一个轻分子,其旋转能级的间距比重转子(如 CO 或 CS)大得多,相应的波长也要短得多(最低跃迁为 0.5 毫米,而 CO 或 CS 为 3-7 毫米)。两个氢原子的核自旋可以是平行的,也可以是反平行的,这导致 H2O 能级分为正交(奇数)和顺磁(偶数)阶梯,其统计权重比分别为 3:1。这两个阶梯之间的辐射跃迁在高阶上是被禁止的,只有在水的一个氢原子与反应物的一个氢原子交换的化学反应中,才能将正交转变为顺磁,反之亦然。

Refer to caption
Figure 1: The near-IR spectrum of the Earth showing the many water vibrational bands together with CO2. The bands below 3 μ\mum are due to overtones and combination bands and are often targeted in exoplanet searches. This spectrum was observed with the NIMS instrument on the Galileo spacecraft during its Earth flyby in December 1990. From Encrenaz (2008), with permission from Annual Reviews, based on Drossart et al. (1993).
图 1:地球的近红外光谱显示了许多水振动带以及 CO 2 。3 μ\mu m 以下的带是由于倍频和组合带引起的,通常在寻找外行星时被瞄准。该光谱是在 1990 年 12 月盖利略号飞船飞越地球期间使用 NIMS 仪器观测到的。根据 Encrenaz(2008)的许可,基于 Drossart 等人(1993)的研究。

Infrared spectroscopy can reveal both water vapor and ice. Water has three active vibrational modes: the fundamental vv=1–0 bands of the ν1\nu_{1} and ν3\nu_{3} symmetric and asymmetric stretches centered at 2.7 μ\mum and 2.65 μ\mum, respectively, and the ν2\nu_{2} bending mode at 6.2 μ\mum. Overtone (Δv=2\Delta v=2 or larger) and combination (e.g., ν2\nu_{2}+ν3\nu_{3}) transitions occur in hot gas at shorter wavelengths (see Fig. 1 for example). Gas-phase water therefore has a rich vibration-rotation spectrum with many individual lines depending on the temperature of the gas. In contrast, the vibrational bands of water ice have no rotational substructure and consist of very broad profiles, with the much stronger ν3\nu_{3} band overwhelming the weak ν1\nu_{1} band. The ice profile shapes depend on the morphology, temperature and environment of the water molecules (Hudgins et al., 1993). Crystalline water ice is readily distinguished by a sharp feature around 3.1 μ\mum that is lacking in amorphous water ice. Libration modes of crystalline water ice are found at 45 and 63 μ\mum (Moore and Hudson, 1994).
红外光谱学可以揭示水蒸气和冰。水有三种活跃振动模式:基本 vv =1-0 带的 ν1subscript1\nu_{1}ν3subscript3\nu_{3} 对称和非对称伸展位于 2.7 μ\mu m 和 2.65 μ\mu m,以及 6.2 μ\mu m 处的 ν2subscript2\nu_{2} 弯曲模式。在较短波长处,热气体中会发生倍频( Δv=22\Delta v=2 或更大)和组合(例如, ν2subscript2\nu_{2} + ν3subscript3\nu_{3} )跃迁(例如,见图 1)。因此,气相水具有丰富的振动-转动光谱,具有许多个体线,取决于气体的温度。相比之下,水冰的振动带没有旋转亚结构,由非常宽的轮廓组成,较强的 ν3subscript3\nu_{3} 带压倒了较弱的 ν1subscript1\nu_{1} 带。冰的轮廓形状取决于水分子的形态、温度和环境(Hudgins 等,1993 年)。结晶水冰可通过缺乏于无定形水冰的 3.1 μ\mu m 附近的尖锐特征轻松区分。结晶水冰的摆动模式位于 45 和 63 μ\mu m 处(Moore 和 Hudson,1994 年)。

Spectra of hydrous silicates (also known as phyllosilicates, layer-lattice silicates or ‘clays’) show sharp features at 2.70–2.75 μ\mum due to isolated OH groups and a broader absorption from 2.75–3.2 μ\mum caused by interlayered (‘bound’) water molecules. At longer wavelengths, various peaks can occur depending on the composition; for example, the hydrous silicate montmorillonite has bands at 49 and 100 μ\mum (Koike et al., 1982).
含水硅酸盐(也称为云母状硅酸盐、层状晶格硅酸盐或“粘土”)的光谱在 2.70-2.75μm 处显示出由孤立的 OH 基团引起的尖锐特征,以及由层间(“结合”)水分子引起的 2.75-3.2μm 处的较宽吸收。在较长波长处,根据组成可以出现各种峰值;例如,含水硅酸盐蒙脱石在 49 和 100μm 处有吸收带(Koike 等,1982 年)。

Bound-bound electronic transitions of water occur at far-UV wavelengths around 1240 Å, but have not yet been detected in space.
水的束缚-束缚电子跃迁发生在远紫外波长约 1240 埃处,但尚未在太空中检测到。

2.3 Water excitation  2.3 水的激发

The strength of an emission or absorption line of water depends on the number of molecules in the telescope beam and, for gaseous water, on the populations of the individual energy levels. These populations, in turn, are determined by the balance between the collisional and radiative excitation and de-excitation of the levels. The radiative processes involve both spontaneous emission and stimulated absorption and emission by a radiation field produced by a nearby star, by warm dust, or by the molecules themselves.
水的发射或吸收线的强度取决于望远镜光束中分子的数量,对于气态水而言,还取决于各个能级的能级的粒子数。这些粒子数又由能级的碰撞激发和辐射激发和去激发之间的平衡决定。辐射过程涉及到自发辐射以及由附近恒星、温暖尘埃或分子本身产生的辐射场的刺激吸收和辐射。

The main collisional partner in interstellar clouds is H2. Accurate state-to-state collisional rate coefficients, CuC_{u\ell}, of H2O with both ortho- and para-H2 over a wide range of temperatures have recently become available thanks to a dedicated chemical physics study (Daniel et al., 2011). Other collision partners such as H, He and electrons are generally less important. In cometary atmospheres, water itself provides most of the collisional excitation.
星际云中的主要碰撞伴侣是 H 2 。最近,由于一项专门的化学物理研究(Daniel 等,2011 年),已经获得了 H 2 O 与正 H 和反 H 2 在广泛温度范围内的准确状态对状态碰撞速率系数 CusubscriptC_{u\ell} 。其他碰撞伴侣如 H、He 和电子通常不那么重要。在彗星大气中,水本身提供了大部分碰撞激发。

Astronomers traditionally analyze molecular observations through a Boltzmann diagram, in which the level populations are plotted versus the energy of the level involved. The slope of the diagram gives the inverse of the excitation temperature. If collisional processes dominate over radiative processes, the populations are in ‘local thermodynamic equilibrium’ (LTE) and the excitation temperature is equal to the kinetic temperature of the gas, Tex=TkinT_{\rm ex}=T_{\rm kin}. Generally level populations are far from LTE and molecules are excited by collisions and de-excited by spontaneous emission, leading to Tex<TkinT_{\rm ex}<T_{\rm kin}. The critical density roughly delineates the transition between these regimes: ncr=Au/Cun_{cr}=A_{u\ell}/C_{u\ell} and therefore scales with μu2νu3\mu_{u\ell}^{2}\nu_{u\ell}^{3}, where AA is the Einstein spontaneous emission coefficient, μ\mu the electric dipole moment and ν\nu the frequency of the transition uu\to\ell. In the case of water, the combination of a large dipole moment (1.86 Debye) and high frequencies results in high critical densities of 10810^{8}10910^{9} cm-3 for pure rotational transitions.
天文学家传统上通过玻尔兹曼图来分析分子观测,其中级别的分布与所涉及级别的能量相对比。图的斜率给出了激发温度的倒数。如果碰撞过程主导辐射过程,那么分布处于‘局部热力学平衡’(LTE)状态,激发温度等于气体的动力学温度。一般来说,级别的分布远离 LTE,分子通过碰撞激发并通过自发辐射去激发,导致。临界密度大致划定了这些状态之间的过渡:,因此与成比例,其中是爱因斯坦自发辐射系数,是电偶极矩,是过渡的频率。在水的情况下,大偶极矩(1.86 Debye)和高频率的组合导致纯旋转跃迁的临界密度很高,为 - 厘米。

Analysis of water lines is much more complex than for simple molecules, such as CO, for a variety of reasons. First, because of the large dipole moment and high frequencies, the rotational transitions of water are usually highly optically thick, even for abundances as low as 101010^{-10}. Second, the water transitions couple effectively with mid- and far-infrared radiation from warm dust, which can pump higher energy levels. Third, the fact that the ‘backbone’ levels with KAK_{A}=0 or 1 have lower radiative decay rates than higher KAK_{A} levels can lead to population ‘inversion’, in which the population in the upper state divided by its statistical weight exceeds that for the lower state (i.e., TexT_{\rm ex} becomes negative). Infrared pumping can also initiate this inversion. The result is the well-known maser phenomenon, which is widely observed in several water transitions in star-forming regions (e.g., Furuya et al., 2003; Neufeld et al., 2013; Hollenbach et al., 2013). The bottom line is that accurate analysis of interstellar water spectra often requires additional independent constraints, for example from H182{}_{2}^{18}O or H172{}_{2}^{17}O isotopologues, whose abundances are reduced by factors of about 550 and 2500, respectively, and whose lines are more optically thin. At infrared wavelengths, lines are often spectrally unresolved, which further hinders the interpretation.
水线的分析比对简单分子(如 CO)要复杂得多,原因有很多。首先,由于水具有较大的偶极矩和高频率,水的旋转跃迁通常具有很高的光学厚度,即使对于丰度低至 1010superscript101010^{-10} 。其次,水的跃迁与来自暖尘埃的中远红外辐射有效耦合,可以激发更高能级。第三,具有 KAsubscriptK_{A} =0 或 1 的‘骨干’能级的辐射衰减速率低于更高 KAsubscriptK_{A} 能级,可能导致人口‘反转’,即上态的人口除以其统计权重超过下态的情况(即 TexsubscriptT_{\rm ex} 变为负值)。红外泵浦也可以引发这种反转。结果就是众所周知的激射现象,在星形成区域的几个水跃迁中广泛观察到(例如,Furuya 等人,2003 年;Neufeld 等人,2013 年;Hollenbach 等人,2013 年)。 准确分析星际水谱往往需要额外的独立约束条件,例如来自 H 182superscriptsubscriptabsent218{}_{2}^{18} O 或 H 172superscriptsubscriptabsent217{}_{2}^{17} O 同位素的约束,它们的丰度分别降低约 550 倍和 2500 倍,其谱线更薄。在红外波长下,谱线通常是未解析的,这进一步阻碍了解释。

Refer to caption
Figure 2: Summary of the main gas-phase and solid-state chemical reactions leading to the formation and destruction of H2O under non-equilibrium conditions. Three different chemical regimes can be distinguished: (i) ion-molecule chemistry, which dominates gas-phase chemistry at low TT; (ii) high-temperature neutral-neutral chemistry; and (iii) solid state chemistry. ee stands for electron, ν\nu for photon and ss-X indicates that species X is on the grains. Simplified version of figure by van Dishoeck et al. (2011).
图 2:在非平衡条件下导致 H 2 O 形成和破坏的主要气相和固态化学反应的总结。可以区分三种不同的化学机制:(i)离子-分子化学,在低 TT 主导气相化学;(ii)高温中性-中性化学;和(iii)固态化学。 ee 代表电子, ν\nu 代表光子, slimit-froms- X 表示物种 X 位于颗粒上。van Dishoeck 等人(2011)的图的简化版本。

2.4 Water chemistry  2.4 水化学

2.4.1 Elemental abundances and equilibrium chemistry
2.4.1 元素丰度和平衡化学

The overall abundance of elemental oxygen with respect to total hydrogen nuclei in the interstellar medium is estimated to be 5.75×1045.75\times 10^{-4} (Przybilla et al., 2008), of which 16–24% is locked up in refractory silicate material in the diffuse interstellar medium (Whittet, 2010). The abundance of volatile oxygen (i.e., not tied up in some refractory form) is measured to be 3.2×1043.2\times 10^{-4} in diffuse clouds (Meyer et al., 1998), so this is the maximum amount of oxygen that can cycle between water vapor and ice in dense clouds. Counting up all the forms of detected oxygen in diffuse clouds, the sum is less than the overall elemental oxygen abudance. Thus, a fraction of oxygen is postulated to be in some yet unknown refractory form, called UDO (‘unknown depleted oxygen’), whose fraction may increase from 20% in diffuse clouds up to 50% in dense star-forming regions (Whittet, 2010). For comparison, the abundances of elemental carbon and nitrogen are 3×1043\times 10^{-4} and 1×1041\times 10^{-4}, respectively, with about 2/3 of the carbon thought to be locked up in solid carbonaceous material.
宇宙间介质中相对于总氢核的元氧总丰度估计为 5.75×1045.75superscript1045.75\times 10^{-4} (Przybilla 等人,2008 年),其中 16-24%锁定在弥散宇宙间介质中的耐火硅酸盐材料中(Whittet,2010 年)。挥发性氧的丰度(即未被某些耐火形式所束缚的氧)在弥散云中被测定为 3.2×1043.2superscript1043.2\times 10^{-4} (Meyer 等人,1998 年),因此这是可以在密集云中水蒸气和冰之间循环的氧的最大量。统计在弥散云中检测到的所有形式的氧,总和小于整体元氧丰度。因此,据推测,一部分氧被假定为某种尚未知的耐火形式,称为 UDO(“未知耗尽氧”),其比例可能从弥散云中的 20%增加到密集星形成区域的 50%(Whittet,2010 年)。作为比较,元碳和氮的丰度分别为 3×1043superscript1043\times 10^{-4}1×1041superscript1041\times 10^{-4} ,其中约有 2/3 的碳被认为锁定在固体碳质材料中。

For a gas in thermodynamic equilibrium (TE), the fractional abundance of water is simply determined by the elemental composition of the gas and the stabilities of the molecules and solids that can be produced from it. For standard interstellar abundances222The notation [X] indicates the overall abundance of element X in all forms, be it atoms, molecules or solids. with [O]/[C]>1>1, there are two molecules in which oxygen can be locked up: CO and H2O. At high pressures in TE, the fraction of CO results from the equilibrium between CO and CH4, with CO favored at higher temperatures. For the volatile elemental abundances quoted above, this results in an H2O fractional abundance of (23)×104(2-3)\times 10^{-4} with respect to total hydrogen, if the CO fractional abundance ranges from (01)×104(0-1)\times 10^{-4}. With respect to H2, the water abundance would then be (56)×104(5-6)\times 10^{-4} assuming that the fraction of hydrogen in atomic form is negligible (the density of hydrogen nuclei nH=nn_{\rm H}=n(H) + 2nn(H2)). Equilibrium chemistry is established at densities above roughly 101310^{13} cm-3, when three body processes become significant. Such conditions are found in planetary atmospheres and in the shielded midplanes of the inner few AU of protoplanetary disks.
在热力学平衡(TE)中,水的分数丰度仅由气体的元素组成和可以从中产生的分子和固体的稳定性决定。对于标准的星际丰度 222The notation [X] indicates the overall abundance of element X in all forms, be it atoms, molecules or solids. ,其中氧可以被锁定在两种分子中:CO 和 H 2 O。在 TE 中的高压下,CO 的分数来自于 CO 和 CH 4 之间的平衡,CO 在较高温度下更有利。对于上述揮发性元素丰度,如果 CO 的分数丰度范围为 (01)×10401superscript104(0-1)\times 10^{-4} ,那么相对于总氢,这将导致 H 2 O 的分数丰度为 (23)×10423superscript104(2-3)\times 10^{-4} 。相对于 H 2 ,水的丰度将是 (56)×10456superscript104(5-6)\times 10^{-4} ,假设氢的原子形式的分数可以忽略不计(氢核的密度 nH=nsubscriptn_{\rm H}=n (H)+ 2 nn (H 2 ))。在密度大约 1013superscript101310^{13} cm -3 以上时,平衡化学在三体过程变得显著时建立。这种条件在行星大气层和原行星盘内几个 AU 的屏蔽中面中发现。

Under most conditions in interstellar space, however, the densities are too low for equilibrium chemistry to be established. Also, strong UV irradiation drives the chemistry out of equilibrium, even in high-density environments, such as the upper atmospheres of planets and disks. Under these conditions, the fractional abundances are determined by the kinetics of the two-body reactions between the various species in the gas. Figure 2 summarizes the three routes to water formation that have been identified. Each of these routes dominates in a specific environment.

2.4.2 Low temperature gas-phase chemistry

In diffuse and translucent interstellar clouds with densities less than 104\sim 10^{4} cm-3 and temperatures below 100 K, water is formed largely by a series of ion-molecule reactions (e.g., Herbst and Klemperer, 1973). The network starts with the reactions O + H+3{}_{3}^{+} and O+ + H2 leading to OH+. The H+3{}_{3}^{+} ion is produced by interactions of energetic cosmic-ray particles with the gas, producing H+2{}_{2}^{+} and H+, with the subsequent fast reaction of H+2{}_{2}^{+} + H2 leading to H+3{}_{3}^{+}. The cosmic ray ionization rate of atomic hydrogen denoted by ζH\zeta_{\rm H} can be as high as 101510^{-15} s-1 in some diffuse clouds, but drops to 101710^{-17} s-1 in denser regions (Indriolo and McCall, 2012; Rimmer et al., 2012). The ionization rate of H2 is ζH22ζH\zeta_{\rm H_{2}}\approx 2\zeta_{\rm H}.
在密度小于 104similar-toabsentsuperscript104\sim 10^{4} cm -3 且温度低于 100 K 的弥散和半透明星际云中,水主要通过一系列离子-分子反应形成(例如 Herbst 和 Klemperer,1973 年)。该网络始于反应 O + H +3superscriptsubscriptabsent3{}_{3}^{+} 和 O + + H 2 导致 OH + 。H +3superscriptsubscriptabsent3{}_{3}^{+} 离子是由高能宇宙射线粒子与气体相互作用产生的,产生 H +2superscriptsubscriptabsent2{}_{2}^{+} 和 H + ,随后 H +2superscriptsubscriptabsent2{}_{2}^{+} + H 2 的快速反应导致 H +3superscriptsubscriptabsent3{}_{3}^{+} 。原子氢的宇宙射线电离率由 ζHsubscript\zeta_{\rm H} 表示,在一些弥散云中可能高达 1015superscript101510^{-15} s -1 ,但在密集区域下降至 1017superscript101710^{-17} s -1 (Indriolo 和 McCall,2012 年;Rimmer 等,2012 年)。H 2 的电离率为 ζH22ζHsubscriptsubscript22subscript\zeta_{\rm H_{2}}\approx 2\zeta_{\rm H}

A series of rapid reactions of OH+ and H2O+ with H2 lead to H3O+, which can dissociatively recombine to form H2O and OH with branching ratios of \sim0.17 and 0.83, respectively (Buhr et al., 2010). H2O is destroyed by photodissociation and by reactions with C+, H+3{}_{3}^{+} and other ions such as HCO+. Photodissociation of H2O starts to be effective shortward of 1800 Å and continues down to the ionization threshold at 983 Å (12.61 eV), including Ly α\alpha at 1216 Å. Its lifetime in the general interstellar radiation field, as given by Draine (1978), is only 40 yr.
OH + 和 H 2 O + 的快速反应系列与 H 2 导致 H 3 O + ,后者可以解离地重新组合形成 H 2 O 和 OH,其分支比分别为 0.17 和 0.83(Buhr 等人,2010)。H 2 O 被光解和与 C + 、H +3superscriptsubscriptabsent3{}_{3}^{+} 和其他离子如 HCO + 的反应所破坏。H 2 O 的光解开始在 1800 Å以下有效,并一直持续到 983 Å(12.61 eV)的电离阈值,包括 1216 Å处的 Ly α\alpha 。根据 Draine(1978)给出的普通星际辐射场中,其寿命仅为 40 年。

2.4.3 High-temperature gas-phase chemistry
2.4.3 高温气相化学

At temperatures above 230 K, the energy barriers for reactions with H2 can be overcome and the reaction O + H2 \to OH + H becomes the dominant channel initiating water formation (Elitzur and Watson, 1978). OH subsequently reacts with H2 to form H2O, a reaction which is exothermic, but has an energy barrier of \sim2100 K (Atkinson et al., 2004). This route drives all the available gas-phase oxygen into H2O, unless strong UV or a high atomic H abundance convert some water back to OH and O. High-temperature chemistry dominates the formation of water in shocks, in the inner envelopes around protostars, and in the warm surface layers of protoplanetary disks.
在高于 230 K 的温度下,与 H 2 反应的能量障碍可以被克服,反应 O + H 2 \to OH + H 成为引发水形成的主要途径(Elitzur 和 Watson,1978 年)。随后 OH 与 H 2 反应形成 H 2 O,这是一个放热反应,但有一个 similar-to\sim 2100 K 的能量障碍(Atkinson 等人,2004 年)。除非强紫外线或高原子 H 丰度将一些水转化回 OH 和 O,否则这条途径将所有可用的气相氧转化为 H 2 O。在冲击中,原恒星周围的内部包层以及原行星盘的温暖表面层中,高温化学主导了水的形成。

2.4.4 Ice chemistry  2.4.4 冰化学

The timescale for an atom or molecule to collide with a grain and stick to it is tfo=3×109/nH2t_{\rm fo}=3\times 10^{9}/n_{\rm H_{2}} yr for normal size grains and sticking probabilities close to unity (Hollenbach et al., 2009). Thus, for densities greater than 10410^{4} cm-3, the time scales for freeze-out are less than a few ×105\times 10^{5} yr, generally smaller than the lifetime of dense cores (at least 10510^{5} yr). Reactions involving dust grains are therefore an integral part of the chemistry. Even weakly bound species, such as atomic H, have a long enough residence time on the grains at temperatures of 10–20 K to react; H2 also participates in some surface reactions, but remains largely in the gas. Tielens and Hagen (1982) postulated that the formation of water from O atoms proceeds through three routes involving hydrogenation of ss-O, ss-O2 and ss-O3, respectively, where ss-X indicates a species on the surface. All three routes have recently been verified and quantified in the laboratory and detailed networks with simulations have been drawn up (see Cuppen et al., 2010; Oba et al., 2012; Lamberts et al., 2013, for summaries).
原文中有一些数字和专有名词,将其保持不变,其他内容翻译如下: 原文翻译:原子或分子与颗粒碰撞并粘附的时间尺度对于正常大小的颗粒和粘附概率接近于单位(Hollenbach 等,2009 年)为 tfo=3×109/nH2subscript3superscript109subscriptsubscript2t_{\rm fo}=3\times 10^{9}/n_{\rm H_{2}} 年。因此,对于密度大于 104superscript10410^{4} cm -3 的情况,冻结的时间尺度小于几年,通常小于密集核心的寿命(至少 105superscript10510^{5} 年)。因此,涉及尘埃颗粒的反应是化学的一个组成部分。即使是弱结合的物种,如原子 H,在温度为 10-20 K 时在颗粒上有足够长的停留时间来发生反应;H 2 也参与了一些表面反应,但主要保持在气态中。Tielens 和 Hagen(1982 年)假设从 O 原子形成水的过程通过涉及 ss -O, ss -O 2ss -O 3 的氢化三条途径进行,其中 ss -X 表示表面上的一种物种。这三条途径最近已在实验室中得到验证和量化,并绘制了详细的网络和模拟(有关摘要,请参见 Cuppen 等,2010 年;Oba 等,2012 年;Lamberts 等,2013 年)。

Water ice formation is in competition with various desorption processes, which limit the ice build-up. At dust temperatures below the thermal sublimation limit, photodesorption is an effective mechanism to get species back to the gas phase, although only a small fraction of the UV absorptions results in desorption of intact H2O molecules (Andersson and van Dishoeck, 2008). The efficiency is about 10310^{-3} per incident photon, as determined through laboratory experiments and theory (Westley et al., 1995; Öberg et al., 2009; Arasa et al., 2010). Only the top few monolayers of the ice contribute. The UV needed to trigger photodesorption can come either from a nearby star, or from the general interstellar radiation field. Deep inside clouds, cosmic rays produce a low level of UV flux, 104\sim 10^{4} photons cm-2 s-1, through interaction with H2 (Prasad and Tarafdar, 1983). Photodesorption via X-rays is judged to be inefficient, although there are large uncertainties in the transfer of heat within a porous aggregate (Najita et al., 2001). UV photodesorption of ice is thought to dominate the production of gaseous water in cold pre-stellar cores, the cold outer envelopes of protostars and the outer parts of protoplanetary disks.
水冰形成与各种脱附过程竞争,这些过程限制了冰的积累。在低于热升华限制的尘埃温度下,光解脱是一种有效的机制,可以将物种送回气相,尽管只有少量紫外吸收导致完整的 H2O 分子脱附(Andersson 和 van Dishoeck,2008 年)。效率约为每个入射光子 1,通过实验室实验和理论确定(Westley 等,1995 年;Öberg 等,2009 年;Arasa 等,2010 年)。只有冰的顶层几个单层起作用。触发光解脱所需的紫外线可以来自附近的恒星,也可以来自普通的星际辐射场。在云层深处,宇宙射线通过与 H 相互作用产生低水平的紫外线通量,每立方厘米 2 个光子每秒 4 个(Prasad 和 Tarafdar,1983 年)。通过 X 射线的光解脱被认为效率低,尽管在多孔聚集体内热传递存在很大的不确定性(Najita 等,2001 年)。 紫外光解冰被认为主导了在冷的原恒星核心、原恒星的冷外层包层和原行星盘的外部产生气态水的过程。

Other non-thermal ice desorption processes include cosmic ray induced spot heating (which works for CO, but is generally not efficient for strongly bound molecules like H2O) and desorption due to the energy liberated by the reaction (called ‘reactive’ or ‘chemical’ desorption). These processes are less well explored than photodesorption, but a recent laboratory study of ss-D + ss-OD \to ss-D2O suggests that as much as 90% of the product can be released into the gas phase (Dulieu et al., 2013). The details of this mechanism, which has not yet been included in models, are not yet understood and may strongly depend on the substrate.
其他非热冰解吸附过程包括宇宙射线诱导的斑点加热(适用于 CO,但通常对于像 H2O 这样结合力强的分子效率不高)和由反应释放的能量导致的解吸附(称为‘反应性’或‘化学’解吸附)。这些过程比光解吸附研究得更少,但最近一项实验室研究表明, ss -D + ss -OD \to ss -D 2 O 的产物多达 90%可能释放到气相中(Dulieu 等人,2013 年)。这种机制的细节尚未包含在模型中,目前尚不清楚,可能在很大程度上取决于基底。

Once the dust temperature rises above \sim100 K (precise value being pressure dependent), H2O ice thermally sublimates on timescales of years, leading to initial gas-phase abundances of H2O as high as the original ice abundances. These simulations use a binding energy of 5600 K for amorphous ice and a slightly higher value of 5770 K for crystalline ice, derived from laboratory experiments (Fraser et al., 2001). Thermal desorption of ices contributes to the gas-phase water abundance in the warm inner protostellar envelopes (‘hot cores’) and inside the snow line in disks.
一旦尘埃温度升至 similar-to\sim 100 K(精确值取决于压力),H 2 O 冰在几年的时间尺度上热力升华,导致 H 2 O 的初始气相丰度高达原始冰的丰度。这些模拟使用 5600 K 的非晶冰结合能和稍高一点的 5770 K 的结晶冰结合能,这些数值来自实验室实验(Fraser 等人,2001 年)。冰的热解吸对暖内部原恒星包层(“热核”)和盘内雪线内的气相水丰度做出贡献。

2.4.5 Water deuteration  2.4.5 水的重氘化

Deuterated water, HDO and D2O, is formed through the same processes as illustrated in Figure 2. There are, however, a number of chemical processes that can enhance the HDO/H2O and D2O/H2O ratios by orders of magnitude compared with the overall [D]/[H] ratio of 2.0×1052.0\times 10^{-5} in the local interstellar medium (Prodanović et al., 2010). A detailed description is given in the chapter by Ceccarelli et al., here only a brief summary is provided.
重氘化水,HDO 和 D 2 O,是通过与图 2 中所示的相同过程形成的。然而,有许多化学过程可以使 HDO/H 2 O 和 D 2 O/H 2 O 比率增加数个数量级,与局部星际介质中的总[D]/[H]比率 2.0×1052.0superscript1052.0\times 10^{-5} 相比(Prodanović等,2010 年)。Ceccarelli 等人在本章中提供了详细描述,这里仅提供简要总结。

In terms of pure gas-phase chemistry, the direct exchange reaction H2O +HD \leftrightarrow HDO H2 is often considered in solar system models ++(Richet et al., 1977). In thermochemical equilibrium this reaction can provide at most a factor of 3 enhancement, and even that may be limited by kinetics (Lécluse and Robert, 1994). The exchange reaction D + OH \to H + OD, which has a barrier of \sim100 K (Sultanov and Balakrishnan, 2004), is particularly effective in high-temperature gas such as present in the inner disk (Thi et al., 2010b). Photodissociation of HDO enhances OD compared with OH by a factor of 2–3, which could be a route to further fractionation.
就纯气相化学而言,直接交换反应 H 2 O + HD \leftrightarrow HDO H 2 在太阳系模型中经常被考虑(Richet 等,1977 年)。在热化学平衡中,这种反应最多可以提供 3 倍的增强,甚至可能受到动力学的限制(Lécluse 和 Robert,1994 年)。交换反应 D + OH \to H + OD,其势垒为 similar-to\sim 100 K(Sultanov 和 Balakrishnan,2004 年),在内部盘中存在的高温气体中特别有效(Thi 等,2010b 年)。HDO 的光解离使 OD 比 OH 增加 2-3 倍,这可能是进一步分馏的途径。

The bulk of the deuterium fractionation in cold clouds comes from gas-grain processes. Tielens (1983) pointed out that the fraction of deuterium relative to hydrogen atoms arriving on a grain surface, D/H, is much higher than the overall [D]/[H] ratio, which can be implanted into molecules in the ice. This naturally leads to enhanced formation of OD, HDO and D2O ice according to the grain-surface formation routes. The high atomic D/H ratio in the gas arises from the enhanced gaseous H2D+, HD+2{}_{2}^{+}, and D+3{}_{3}^{+} abundances at low temperatures (\leq25 K), when the ortho-H2 abundance drops and their main destroyer, CO, freezes out on the grains (Pagani et al., 1992; Roberts et al., 2003). Dissociative recombination with electrons then produces enhanced D. The enhanced H2D+ also leads to enhanced H2DO+ and thus HDO in cold gas, but this is usually a minor route compared with gas-grain processes.
冷云中重水素分馏的大部分来自气相-颗粒过程。Tielens(1983)指出,相对于氢原子到达颗粒表面的重水素分数,D/H,要比整体[D]/[H]比例高得多,这可以被植入到冰中的分子。这自然导致了根据颗粒表面形成途径增强 OD、HDO 和 D0O 冰的形成。气体中的高原子 D/H 比率来自于低温(25K 以下)时增强的气态 H1D2、HD3 和 D4 的丰度,当正氢的丰度下降并且它们的主要破坏者 CO 在颗粒上冻结时(Pagani 等,1992; Roberts 等,2003)。随后与电子的解离复合产生增强的 D。增强的 H7D8 也导致了冷气中增强的 H9DO10 和因此 HDO,但与气相-颗粒过程相比,这通常是一个次要途径。

On the grains, tunneling reactions can have the opposite effect, reducing the deuterium fractionation. For example, the OD + H2 tunneling reaction producing HDO ice is expected to occur slower than the OH + H2 reaction leading to H2O ice. On the other hand, thermal exchange reactions in the ice, such as H2O + OD \to HDO + OH have been shown to occur rapidly in ices at higher temperatures; these can both enhance and decrease the fractionation. Both thermal desorption at high ice temperatures and photodesorption at low ice temperatures have a negligible effect on the deuterium fractionation, i.e., the gaseous HDO/H2O and D2O/H2O ratios reflect the ice ratios if no other gas-phase processes are involved.
在颗粒上,隧道反应可能产生相反的效果,降低氘分馏。例如,OD + H 的隧道反应产生 HDO 冰的速度预计比 OH + H 反应产生 H2O 冰的速度慢。另一方面,在冰中的热交换反应,如 H3O + OD → HDO + OH,已被证明在较高温度下的冰中迅速发生;这些反应既可以增强又可以减少分馏。在高冰温度下的热解吸和在低冰温度下的光解吸对氘分馏几乎没有影响,即气态 HDO/H2O 和 DO/H2O 比例反映了冰的比例,如果没有其他气相过程参与。

3 CLOUDS AND PRE-STELLAR CORES: ONSET OF WATER FORMATION
3 云和原恒星核:水形成的开始

In this and following sections, our knowledge of the water reservoirs during the various evolutionary stages from clouds to planets will be discussed. The focus is on low-mass protostars (100 L) and pre-main sequence stars (spectral type A or later). Unless stated otherwise, fractional abundances are quoted with respect to H2 and are simply called ‘abundances’. Often the denominator, i.e., the (column) density of H2, is more uncertain than the numerator.<<
在本节和接下来的部分中,我们将讨论从云到行星各个演化阶段的水库知识。重点放在低质量原恒星(100 L )和主序星前期(A 型或更晚的光谱类型)。除非另有说明,分数丰度都是相对于 H 2 并简称为“丰度”。通常分母,即 H 2 的(柱)密度,比分子更不确定。 <<

The bulk of the water in space is formed on the surfaces of dust grains in dense molecular clouds. Although a small amount of water is produced in the gas in diffuse molecular clouds through ion-molecule chemistry, its abundance of 108\sim 10^{-8} found by Herschel-HIFI (Flagey et al., 2013) is negligible compared with that produced in the solid state. In contrast, observations of the 3 μ\mum water ice band toward numerous infrared sources behind molecular clouds, from the ground and from space, show that water ice formation starts at a threshold extinction of AV3A_{V}\approx 3 mag (Whittet et al., 2013). These clouds have densities of at least 1000 cm-3, but are not yet collapsing to form stars. The ice abundance is ss-H2O/H25×105{}_{2}\approx 5\times 10^{-5}, indicating that a significant fraction of the available oxygen has been transformed to water ice even at this early stage (Whittet et al., 1988; Murakawa et al., 2000; Boogert et al., 2011). Such high ice abundances are too large to result from freeze-out of gas-phase water produced by ion-molecule reactions.
宇宙中大部分水是在密集的分子云尘埃颗粒表面形成的。尽管少量水是通过离子-分子化学在弥散的分子云气体中产生的,但由 Herschel-HIFI(Flagey 等人,2013 年)发现的其丰度与固态产生的水相比微不足道。相比之下,对分子云后面众多红外源的 3 μ\mu m 水冰带的观测,无论是从地面还是从太空,都表明水冰形成始于 AV3subscript3A_{V}\approx 3 mag 的阈值消光(Whittet 等人,2013 年)。这些云的密度至少为 1000 cm -3 ,但尚未坍缩形成恒星。冰的丰度为 ss -H 2 O/H 25×105{}_{2}\approx 5\times 10^{-5} ,表明即使在这个早期阶段(Whittet 等人,1988 年;Murakawa 等人,2000 年;Boogert 等人,2011 年),可用氧的相当一部分已经转化为水冰。这种高冰的丰度太大,无法由离子-分子反应产生的气相水冷凝而来。

Refer to caption
Figure 3: Herschel-HIFI spectra of the H2O 1101_{10}1011_{01} line at 557 GHz in a pre-stellar core (top), protostellar envelope (middle) and two protoplanetary disks (bottom) (spectra shifted vertically for clarity). The red dashed line indicates the rest velocity of the source. Note the different scales: water vapor emission is strong toward protostars, but very weak in cold cores and disks. The feature at -15 km s-1 in the TW Hya spectrum is due to NH3. Figure by L. Kristensen, adapted from Caselli et al. (2012), Kristensen et al. (2012) and Hogerheijde et al. (2011, and in prep.).
图 3:在一个原恒星核心(顶部)、原恒星包层(中部)和两个原行星盘(底部)中,557 GHz 处 H 2 O 110subscript1101_{10}101subscript1011_{01} 线的 Herschel-HIFI 光谱(为了清晰起见,光谱在垂直方向上进行了移动)。红色虚线表示源的静止速度。请注意不同的比例尺:水蒸气排放朝向原恒星很强,但在冷核心和盘中非常微弱。在 TW Hya 光谱中-15 km s -1 处的特征是由 NH 3 引起的。图由 L. Kristensen 绘制,改编自 Caselli 等人(2012 年)、Kristensen 等人(2012 年)和 Hogerheijde 等人(2011 年,准备中)。

The densest cold cores just prior to collapse have such high extinctions that direct IR ice observations are not possible. In contrast, the water reservoir (gas plus ice) can be inferred from Herschel-HIFI observations of such cores. Fig. 3 presents the detection of the H2O 1101_{10}1011_{01} 557 GHz line toward L1544 (Caselli et al., 2012). The line shows blue-shifted emission and red-shifted absorption, indicative of inward motions in the core. Because of the high critical density of water, the emission indicates that water vapor must be present in the dense central part. The infalling red-shifted gas originates on the near-side. Because the different parts of the line profile probe different parts of the core, the line shape can be used to reconstruct the water vapor abundance as a function of position throughout the entire core.
在坍缩前最密集的冷核心具有如此高的消光,以至于无法直接观测到红外冰。相比之下,可以通过 Herschel-HIFI 对这些核心的观测来推断水库(气体加冰)的存在。图 3 展示了朝向 L1544(Caselli 等人,2012 年)的 H 2 O 110subscript1101_{10} - 101subscript1011_{01} 557 GHz 线的检测。该线显示出蓝移发射和红移吸收,表明核心内部存在向内运动。由于水的临界密度很高,发射表明水蒸气必须存在于密集的中心部分。向内运动的红移气体起源于近侧。由于线谱的不同部分探测核心的不同部分,线形状可用于重建整个核心位置处水蒸气丰度的分布。

The best-fit water abundance profile is obtained with a simple gas-grain model, in which atomic O is converted into water ice on the grains, with only a small fraction returned back into the gas by photodesorption (Bergin et al., 2000; Roberts and Herbst, 2002; Hollenbach et al., 2009). The maximum gas-phase water abundance of 107\sim 10^{-7} occurs in a ring at the edge of the core around AV4A_{V}\approx 4 mag, where external UV photons can still penetrate to photodesorb the ice, but where they are no longer effective in photodissociating the water vapor. In the central shielded part of the core, cosmic ray induced UV photons keep a small, 109\sim 10^{-9}, but measurable fraction of water in the gas (Caselli et al., 2012). Quantitatively, the models indicate that the bulk of the available oxygen has been transformed into water ice in the core, with an ice abundance of 104\sim 10^{-4} with respect to H2.
通过一个简单的气体-颗粒模型获得了最佳拟合的水丰度剖面,在该模型中,原子氧在颗粒上转化为水冰,只有一小部分通过光解吸附返回到气相中(Bergin 等,2000 年;Roberts 和 Herbst,2002 年;Hollenbach 等,2009 年)。气相水的最大丰度 107similar-toabsentsuperscript107\sim 10^{-7} 出现在核心边缘的一个环上,约 AV4subscript4A_{V}\approx 4 磁通量处,外部紫外光子仍然可以穿透以光解吸附冰,但在那里它们不再有效地光解水蒸气。在核心的中央受屏蔽部分,宇宙射线诱导的紫外光子保持了一小部分水在气相中(Caselli 等,2012 年)。定量上,模型表明大部分可用氧已经转化为核心中的水冰,其冰丰度相对于 H 为 104similar-toabsentsuperscript104\sim 10^{-4}

4 PROTOSTARS AND OUTFLOWS
4 原恒星和流出

4.1 Outflows  4.1 流出

Herschel-HIFI and PACS data show strong and broad water profiles characteristic of shocks associated with embedded protostars, from low to high mass. In fact, for low-mass protostars this shocked water emission completely overwhelms the narrower lines from the bulk of the collapsing envelope, even though the shocks contain less than 1% of the mass of the system. Maps of the water emission around solar-mass protostars such as L1157 reveal water not only at the protostellar position but also along the outflow at ‘hot spots’ where the precessing jet interacts with the cloud (Nisini et al., 2010). Thus, water traces the currently shocked gas at positions, which are somewhat offset from the bulk of the cooler entrained outflow gas seen in the red- and blue-shifted lobes of low-JJ CO lines (Tafalla et al., 2013; Lefloch et al., 2010).
赫歇尔-HIFI 和 PACS 数据显示出与嵌入式原恒星相关的冲击特征强烈而宽广的水谱线,从低质量到高质量。实际上,对于低质量原恒星,这种受到冲击的水发射完全压倒了来自坍缩包层大部分的较窄线,即使这些冲击所含的质量不到系统总质量的 1%。围绕太阳质量原恒星的水发射图,如 L1157,显示出水不仅存在于原恒星位置,还存在于与云层相互作用的“热点”处的流出口(Nisini 等人,2010 年)。因此,水迹象目前受到冲击的气体在位置上的存在,这些位置与在低 JJ CO 线的红移和蓝移叶瓣中看到的较冷流出气体的大部分有些偏移(Tafalla 等人,2013 年;Lefloch 等人,2010 年)。

Determinations of the water abundance in shocks vary from values as low as 10710^{-7} to as high as 10410^{-4} (see van Dishoeck et al., 2013, for summary). In non-dissociative shocks, the temperature reaches values of a few thousand K and all available oxygen is expected to be driven into water (Kaufman and Neufeld, 1996). The low values likely point to the importance of UV radiation in the shock chemistry and shock structure. For the purposes of this chapter, the main point is that even though water is rapidly produced in shocks at potentially high abundances, the amount of water contained in the shocks is small, and, moreover, most of it is lost to space through outflows.
在冲击中确定的水丰度值从低至 107superscript10710^{-7} 到高至 104superscript10410^{-4} 不等(见 van Dishoeck 等人,2013 年,摘要)。在非解离性冲击中,温度达到几千 K 的数值,所有可用的氧都预计会转化为水(Kaufman 和 Neufeld,1996 年)。低值可能指向 UV 辐射在冲击化学和冲击结构中的重要性。对于本章的目的,主要观点是,即使水在冲击中迅速产生,且可能丰度很高,但冲击中含有的水量很少,而且大部分通过流出失去到太空中。

Refer to caption
Figure 4: Schematic representation of a protostellar envelope and embedded disk with key steps in the water chemistry indicated. Water ice is formed in the parent cloud before collapse and stays mostly as ice until the ice sublimation temperature of \sim100 K close to the protostar is reached. Hot water is formed in high abundances in shocks associated with the outflow, but this water is not incorporated into the planet-forming disk. Figure by R. Visser, adapted from Herbst and van Dishoeck (2009).
图 4:原恒星包层和嵌入式盘的示意图,显示了水化学的关键步骤。水冰在母云中形成,直到接近原恒星的 100K 的冰升华温度为止,大部分仍保持为冰。在与流出相关的冲击中形成高丰度的热水,但这种水不会被纳入形成行星的盘中。图片由 R. Visser 制作,改编自 Herbst 和 van Dishoeck(2009 年)。

4.2 Protostellar envelopes: the cold outer reservoir
4.2 原恒星包层:冷外部储库

As the cloud collapses to form a protostar in the center, the water-ice coated grains created in the natal molecular cloud move inward, feeding the growing star and its surrounding disk (Fig. 4). The water ice abundance can be measured directly through infrared spectroscopy of various water ice bands toward the protostar itself. Close to a hundred sources have been observed, from very low luminosity objects (‘proto-brown dwarfs’) to the highest mass protostars (Gibb et al., 2004; Pontoppidan et al., 2004; Boogert et al., 2008; Zasowski et al., 2009; Öberg et al., 2011). Inferred ice abundances with respect to H2 integrated along the line of sight are (0.5–1)×\times10410^{-4}.
当云坍缩形成中心的原恒星时,出生分子云中形成的水冰包覆颗粒向内移动,为不断增长的恒星及其周围的盘提供营养(图 4)。水冰丰度可以通过直接测量朝向原恒星本身的各种水冰带的红外光谱来确定。已观测到近百个源,从非常低亮度的物体(“原棕矮星”)到最高质量的原恒星(Gibb 等,2004 年;Pontoppidan 等,2004 年;Boogert 等,2008 年;Zasowski 等,2009 年;Öberg 等,2011 年)。沿视线积分的相对于 H 2 的推断冰丰度为(0.5-1) ×\times 104superscript10410^{-4}

The water vapor abundance in protostellar envelopes is probed through spectrally-resolved Herschel-HIFI lines. Because the gaseous water line profiles are dominated by broad outflow emission (Fig. 3), this component needs to be subtracted, or an optically thin water isotopologue needs to be used to determine the quiescent water. Clues to the water vapor abundance structure can be obtained through narrow absorption and emission features in so-called (inverse) P-Cygni profiles (see NGC 1333 IRAS 4A in Fig. 3). The analysis of these data proceeds along the same lines as for pre-stellar cores. The main difference is that the dust temperature now increases inwards, from a low value of 10–20 K at the edge to a high value of several hundred K in the center of the core (Fig. 4). In the simplest spherically symmetric case, the density follows a power-law nrpn\propto r^{-p} with pp=1–2. As for pre-stellar cores, the data require the presence of a photodesorption layer at the edge of the core with a decreasing water abundance at smaller radii, where gaseous water is maintained by the cosmic ray induced photodesorption of water ice (Coutens et al., 2012; Mottram et al., 2013). Analysis of the combined gaseous water and water ice data for the same source shows that the ice/gas ratio is at least 10410^{4} (Boonman and van Dishoeck, 2003). Thus, the bulk of the water stays in the ice in this cold part, at a high abundance of 104\sim 10^{-4} as indicated by direct measurements of both the water ice and gas.
通过光谱分辨的赫歇尔-高保真光谱线 (Herschel-HIFI) 探测原恒星包层中的水蒸气丰度。由于气态水谱线轮廓主要由宽流出发射(图 3)构成,因此需要减去该成分,或者使用光学薄的水同位素体来确定静止水。水蒸气丰度结构的线索可以通过所谓的(逆)天鹅座 P 星云轮廓(见图 3 中的 NGC 1333 IRAS 4A)中较窄的吸收和发射特征获得。这些数据的分析与前恒星核的分析方法相同。主要区别在于,尘埃温度现在向内升高,从边缘的 10-20 K 的低值到核心中心的几百 K 的高值(图 4)。在最简单的球对称情况下,密度遵循幂律 nrpproportional-tosuperscriptn\propto r^{-p} ,其中 pp =1-2。对于恒星前核心,数据要求在核心边缘存在光解吸层,在较小的半径处水丰度会减少,其中气态水由宇宙射线诱导的水冰光解吸来维持(Coutens 等人,2012 年;Mottram 等人,2013 年)。 对同一源的混合气态水和水冰数据的分析显示,冰/气体比至少为 104superscript10410^{4} (Boonman 和 van Dishoeck,2003)。因此,在这个寒冷部分,大部分水保留在冰中,其丰度高达 104similar-toabsentsuperscript104\sim 10^{-4} ,如直接测量的水冰和气体所示。

4.3 Protostellar envelopes: the warm inner part
4.3 原恒星包层:温暖的内部部分

When the infalling parcel enters the radius at which the dust temperature reaches \sim100 K, the gaseous water abundance jumps from a low value around 101010^{-10} to values as high as 10410^{-4} (e.g., Boonman et al., 2003; Herpin et al., 2012; Coutens et al., 2012). The 100 K radius scales roughly as 2.3×1014(L/L)2.3\times 10^{14}\sqrt{(}L/L_{\odot}) cm (Bisschop et al., 2007), and is small, 100 AU, for low-mass sources and a few thousand AU for high-mass protostars. The precise abundance of water in the warm gas is still uncertain, however, and can range from <<10610^{-6}10410^{-4} depending on the source and analysis (Emprechtinger et al., 2013; Visser et al., 2013). A high water abundance would indicate that all water sublimates from the grains in the ‘hot core’ before the material enters the disk; a low abundance the opposite.
当下落的包裹进入尘埃温度达到 similar-to\sim 100 K 的半径时,气态水的丰度从约 1010superscript101010^{-10} 的低值跳升至高达 104superscript10410^{-4} 的值(例如,Boonman 等,2003 年;Herpin 等,2012 年;Coutens 等,2012 年)。100 K 半径大致按 2.3×1014(L/L)2.3\times 10^{14}\sqrt{(}L/L_{\odot}) 厘米缩放(Bisschop 等,2007 年),对于低质量源为 100 AU,对于高质量原恒星为几千 AU。然而,温暖气体中水的确切丰度仍然不确定,可以在 << 106superscript10610^{-6}104superscript10410^{-4} 之间变化,取决于源和分析(Emprechtinger 等,2013 年;Visser 等,2013 年)。高水丰度将表明所有水在物质进入盘之前从颗粒中升华在“热核心”中;低丰度则相反。

The fate of water in protostellar envelopes on scales of the size of the embedded disk is currently not well understood, yet it is a crucial step in the water trail from clouds to disks. To probe the inner few hundred AU, a high excitation line of a water isotopolog line not dominated by the outflow or high angular resolution is needed: ground-based millimeter interferometry of the H182{}_{2}^{18}O 3132203_{13}-2_{20} (Eu=204E_{u}=204 K) line at 203 GHz (Persson et al., 2012) and deep Herschel-HIFI spectra of excited H182{}_{2}^{18}O or H172{}_{2}^{17}O lines, such as the 3123033_{12}-3_{03} (Eu=249E_{u}=249 K) line at 1095 GHz have been used. Two main problems need to be faced in the analysis. First, comparison of ground-based and Herschel lines for the same source show that the high frequency HIFI lines can be optically thick even for H218O and H172{}_{2}^{17}O, because of their much higher Einstein AA coefficients. Second, the physical structure of the envelope and embedded disk on scales of a few hundred AU is not well understood (Jørgensen et al., 2005), so that abundances are difficult to determine since the column of warm H2 is poorly constrained. Compact flattened dust structures are not necessarily disks in Keplerian rotation (Chiang et al., 2008) and only a fraction of this material may be at high temperatures.
目前,在嵌入盘状结构尺度的原恒星包层中,水的命运尚不十分清楚,但它是水从云层到盘状结构轨迹中的关键一步。为了探测内部几百个天文单位,需要一条不受外流或高角分辨率支配的水同位素谱线的高激发线:已经使用了地面毫米波干涉测量 203 GHz 的 H 182superscriptsubscriptabsent218{}_{2}^{18} O 313220subscript313subscript2203_{13}-2_{20} ( Eu=204subscript204E_{u}=204 K) 线 (Persson 等人,2012) 和激发 H 182superscriptsubscriptabsent218{}_{2}^{18} O 或 H 172superscriptsubscriptabsent217{}_{2}^{17} O 线的深 Herschel-HIFI 谱,例如 1095 GHz 的 312303subscript312subscript3033_{12}-3_{03} ( Eu=249subscript249E_{u}=249 K) 线。分析中需要面对两个主要问题。首先,通过对比同一源的地面谱线和赫歇尔谱线,我们发现即使是 H 2 18 O 和 H 172superscriptsubscriptabsent217{}_{2}^{17} O,其高频 HIFI 谱线也可能具有光学厚度,因为它们的爱因斯坦 AA 系数要高得多。其次,几百天文单位尺度的包层和嵌入盘的物理结构尚不清楚(Jørgensen et al., 2005),因此由于暖 H 2 的约束条件较差,其丰度难以确定。 紧凑扁平的尘埃结构不一定是开普勒旋转中的盘状结构(Chiang 等人,2008 年),而且这些材料中只有一小部分可能处于高温状态。

Jørgensen and van Dishoeck (2010b) and Persson et al. (2012) measure water columns and use H2 columns derived from continuum interferometry data on the same scales (\sim1′′) to determine water abundances of 108105\sim 10^{-8}-10^{-5} for three low-mass protostars, consistent with the fact that the bulk of the gas on these scales is cold and water is frozen. From a combined analysis of the interferometric and HIFI data, using C18O 9–8 and 10–9 data to determine the warm H2 column, Visser et al. (2013) infer water abundances of 2×1052×1042\times 10^{-5}-2\times 10^{-4} in the \geq100 K gas, as expected for the larger-scale hot cores.
Jørgensen 和 van Dishoeck(2010b)以及 Persson 等人(2012 年)测量水柱,并使用从相同尺度上的连续干涉数据导出的 H 2 柱密度( similar-to\sim 1 ′′ )来确定三颗低质量原恒星的水丰度,这与这些尺度上大部分气体是冷的且水结冰的事实一致。通过对干涉和 HIFI 数据的联合分析,使用 C 18 O 9-8 和 10-9 数据来确定温暖的 H 2 柱密度,Visser 等人(2013 年)推断出 100K 气体中的水丰度为 2×1052×1042superscript1052superscript1042\times 10^{-5}-2\times 10^{-4} ,这与更大尺度热核心的预期相符。

The important implication of these results is that the bulk of the water stays as ice in the inner few hundred AU and that only a few % of the dust may be at high enough temperatures to thermally sublimate H2O . This small fraction of gas passing through high-temperature conditions for ice sublimation is consistent with 2D models of collapsing envelope and disk formation, which give fractions of <120<1-20% depending on initial conditions (Visser et al., 2009, 2011; Ilee et al., 2011; Harsono et al., 2013; Hincelin et al., 2013).
这些结果的重要含义是,大部分水保持为冰在内部几百 AU,只有少数%的尘埃可能温度足够高以使 H2O 热力升华。这小部分气体通过高温条件进行冰升华与 2D 模型的坍缩包层和盘形成一致,这些模型根据初始条件给出不同的百分比,取决于(Visser 等人,2009 年,2011 年;Ilee 等人,2011 年;Harsono 等人,2013 年;Hincelin 等人,2013 年)。

4.4 Entering the disk: the accretion shock and history of water in disks
进入盘中:吸积冲击和盘中水的历史

The fact that only a small fraction of the material within a few hundred AU radius is at
在几百 AU 半径范围内,只有很小一部分物质处于 100K(§4.3),这意味着大部分水以冰的形式存在,并仍在向内运动(图 4)。然而,在某个半径处,高速下降的包必须遇到低速嵌入式盘,导致边界处的冲击。这种冲击导致冲击前方的尘埃温度比恒星加热所达到的温度更高(参见;参见简单的拟合公式),也可以溅射冰。在早期,吸积速度很高,所有冰都会升华或经历足够强的冲击以诱导溅射。然而,这些物质通常最终进入恒星而不是盘中,因此对当前故事并不感兴趣。人们认为盘的大部分是通过后来坍缩过程中下落的包层吸积而成的,这些包层在大半径进入盘中,冲击要弱得多。
\geq 100 K (§ 4.3) implies that most of the water is present as ice and is still moving inwards (Fig. 4). At some radius, however, the high-velocity infalling parcels must encounter the low-velocity embedded disk, resulting in a shock at the boundary. This shock results in higher dust temperatures behind the shock front than those achieved by stellar heating (
在几百 AU 半径范围内,只有很小一部分物质处于 100K(§4.3),这意味着大部分水以冰的形式存在,并仍在向内运动(图 4)。然而,在某个半径处,高速下降的包必须遇到低速嵌入式盘,导致边界处的冲击。这种冲击导致冲击前方的尘埃温度比恒星加热所达到的温度更高(参见;参见简单的拟合公式),也可以溅射冰。在早期,吸积速度很高,所有冰都会升华或经历足够强的冲击以诱导溅射。然而,这些物质通常最终进入恒星而不是盘中,因此对当前故事并不感兴趣。人们认为盘的大部分是通过后来坍缩过程中下落的包层吸积而成的,这些包层在大半径进入盘中,冲击要弱得多。
Neufeld and Hollenbach 1994; see   参见Visser et al. 2009 for a simple fitting formula) and can also sputter ices. At early times, accretion velocities are high and all ices would sublimate or experience a shock strong enough to induce sputtering. However, this material normally ends up in the star rather than in the disk, so it is not of interest for the current story. The bulk of the disk is thought to be made up through layered accretion of parcels that fall in later in the collapse process, and which enter the disk at large radii, where the shock is much weaker
用于简单拟合公式的更高尘埃温度,也可以溅射冰。在早期,吸积速度很高,所有冰都会升华或经历足够强的冲击以诱导溅射。然而,这些物质通常最终进入恒星而不是盘中,因此对当前故事并不感兴趣。人们认为盘的大部分是通过后来坍缩过程中下落的包层吸积而成的,这些包层在大半径进入盘中,冲击要弱得多。
(Visser et al., 2009)  (Visser 等,2009 年). Indeed, the narrow line widths of H
。事实上,H 的窄线宽
182{}_{2}^{18}O of only 1 km s-1 seen in the interferometric data
O 仅为 1 km s -1 在干涉数据中可见
(Jørgensen and van Dishoeck, 2010b)
(Jørgensen 和 van Dishoeck,2010b)
argue against earlier suggestions, based on Spitzer data, of large amounts of hot water going through an accretion shock in the embedded phase, or even being created through high-temperature chemistry in such a shock
根据斯皮策数据反驳了早期关于大量热水通过内嵌相中的凝聚冲击,甚至是通过这种冲击中的高温化学反应而产生的建议
(Watson et al., 2007)  (Watson 等人,2007). This view that accretion shocks do not play a role also contrasts with the traditional view in the solar system community that all ices evaporate and recondense when entering the disk
。这种认为凝聚冲击不起作用的观点也与太阳系社区中的传统观点相矛盾,即所有冰进入盘时都会蒸发和重新凝结
(Lunine et al., 1991; Owen and Bar-Nun, 1993)
(Lunine 等人,1991;Owen 和 Bar-Nun,1993)
.

Refer to caption
Figure 5: Schematic view of the history of H2O gas and ice throughout a young disk at the end of the accretion phase. The main oxygen reservoir is indicated for each zone. The percentages indicate the fraction of disk mass contained in each zone. Zone 1 contains pristine H2O formed prior to star formation and never altered during the trajectory from cloud to disk. In Zone 7, the ice has sublimated once and recondensed again. Thus, the ice in planet- and planetesimal-forming zones of disks is a mix of pristine and processed ice. From Visser et al. (2011).
图 5:在吸积阶段结束时,年轻盘中 H 2 O 气体和冰的历史的示意图。每个区域都标明了主要的氧储库。百分比表示每个区域中包含的盘质量的比例。第 1 区域包含在恒星形成之前形成且在从云到盘的轨迹中从未改变的原始 H 2 O。在第 7 区域,冰曾经升华过一次并再次凝结。因此,盘中行星和小行星形成区域的冰是原始和经过加工的冰的混合物。来自 Visser 等人(2011 年)。

Figure 5 shows the history of water molecules in disks at the end of the collapse phase at tacc=2.5×105t_{\rm acc}=2.5\times 10^{5} yr for a standard model with an initial core mass of 1 M, angular momentum Ω0=1014\Omega_{0}=10^{-14} s-1 and sound speed cs=0.26c_{s}=0.26 km s-1 (Visser et al., 2011). The material ending up in zone 1 is the only water that is completely ‘pristine’, i.e., formed as ice in the cloud and never sublimated, ending up intact in the disk. Material ending up in the other zones contains water that sublimated at some point along the infalling trajectory. In zones 2, 3 and 4, close to the outflow cavity, most of the oxygen is in atomic form due to photodissociation, with varying degrees of subsequent reformation. In zones 5 and 6, most oxygen is in gaseous water. Material in zone 7 enters the disk early and comes close enough to the star to sublimate. This material does not end up in the star, however, but is transported outward in the disk to conserve angular momentum, re-freezing when the temperature becomes low enough. The detailed chemistry and fractions of water in each of these zones depend on the adopted physical model and on whether vertical mixing is included (Semenov and Wiebe, 2011), but the overall picture is robust.
图 5 显示了在标准模型中,初始核心质量为 1 M ,角动量 Ω0=1014subscript0superscript1014\Omega_{0}=10^{-14} s -1 和声速 cs=0.26subscript0.26c_{s}=0.26 km s -1 (Visser 等,2011 年)的情况下,在坍缩阶段结束时盘中水分子的历史, tacc=2.5×105subscript2.5superscript105t_{\rm acc}=2.5\times 10^{5} 年。最终进入区域 1 的物质是唯一完全“原始”的水,即在云中形成并从未升华的冰,最终完整地进入盘中。进入其他区域的物质包含在某个时刻沿着向内运动轨迹升华的水。在 2、3 和 4 区,靠近流出腔的地方,大部分氧以原子形式存在,由于光解而有不同程度的随后再形成。在 5 和 6 区,大部分氧以气态水形式存在。进入区域 7 的物质早期进入盘中,并且靠近恒星足够接近以升华。然而,这种物质并未最终进入恒星,而是在盘中向外运输以保持角动量,在温度足够低时重新冻结。 这些区域中水的详细化学和分数取决于采用的物理模型以及是否包括垂直混合(Semenov 和 Wiebe,2011 年),但总体情况是稳健的。

5 PROTOPLANETARY DISKS  5 原行星盘

Once accretion stops and the envelope has dissipated, a pre-main sequence star is left, surrounded by a disk of gas and dust. These protoplanetary disks form the crucial link between material in clouds and that in planetary systems. Thanks to the new observational facilities, combined with sophisticated disk chemistry models, the various water reservoirs in disks are now starting to be mapped out. Throughout this chapter, we will call the disk out of which our own solar system formed the ‘solar nebula disk’. 333Alternative nomenclatures in the literature include ‘primordial disk’, ‘presolar disk’, ‘protosolar nebula’ or ‘primitive solar nebula’.
一旦吸积停止并且包层消散,就会留下一个前主序星,周围环绕着一圈气体和尘埃的盘。这些原行星盘形成了云中物质与行星系统中物质之间的关键联系。由于新的观测设施,结合复杂的盘化学模型,现在开始绘制盘中各种水库。在本章中,我们将称我们太阳系形成的盘为“太阳星云盘”。 333Alternative nomenclatures in the literature include ‘primordial disk’, ‘presolar disk’, ‘protosolar nebula’ or ‘primitive solar nebula’.

5.1 Hot and cold water in disks: observations
5.1 盘中的热水和冷水:观测

With increasing wavelengths, regions further out and deeper into the disk can be probed. The surface layers of the inner few AU of disks are probed by near- and mid-IR observations. Spitzer-IRS detected a surprising wealth of highly-excited pure rotational lines of warm water at 10–30
随着波长增加,可以探测到更远和更深处的区域。通过近红外和中红外观测可以探测到盘的内部几 AU 的表面层。Spitzer-IRS 探测到了大量高度激发的暖水纯旋转线,在 10-30
μ\mum (Carr and Najita, 2008; Salyk et al., 2008)
(Carr 和 Najita,2008; Salyk 等,2008)
, and these lines have since been shown to be ubiquitous in disks around low-mass T Tauri stars
,这些线后来被证明在低质量 T Tauri 恒星周围的盘中是普遍存在的
(Pontoppidan et al., 2010a; Salyk et al., 2011)
(Pontoppidan 等,2010a; Salyk 等,2011)
, with line profiles consistent with a disk origin
,具有与盘源一致的线型剖面
(Pontoppidan et al., 2010b)
(Pontoppidan 等人,2010b)
. Typical water excitation temperatures are
。典型的水激发温度为
TexT_{\rm ex}\approx450 K. Spectrally resolved ground-based near-IR vibration-rotation lines around 3
450 K。在地面附近的近红外振动-转动线约 3 的光谱分辨率。
μ\mum show that in some sources the water originates in both a disk and a slow disk wind
显示一些来源的水起源于盘和缓慢的盘风
(Salyk et al., 2008; Mandell et al., 2012)
(Salyk 等,2008; Mandell 等,2012)
. Abundance ratios are difficult to extract from the observations, because the lines are highly saturated and, in the case of Spitzer data, spectrally unresolved. Also, the IR lines only probe down to moderate height in the disk until the dust becomes optically thick. Nevertheless, within the more than an order of magnitude uncertainty, abundance ratios of H2O/CO
丰度比很难从观测数据中提取,因为谱线高度饱和,在 Spitzer 数据的情况下,谱线无法分辨。此外,红外线谱线只能探测到盘中的中等高度,直到尘埃变得光学厚。尽管如此,在超过一个数量级的不确定性范围内,已经推断出了 H 2 O/CO 的丰度比为 1-10,对应的辐射半径高达几个天文单位
\sim1–10 have been inferred for emitting radii up to a few AU
1-10 已被推断出用于发射半径高达几个天文单位
(Salyk et al., 2011; Mandell et al., 2012)
(Salyk 等人,2011 年; Mandell 等人,2012 年)
. This indicates that the inner disks have high water abundances of order
这表明内部盘具有高水含量的数量级
104\sim 10^{-4} and are thus not dry, at least not in their surface layers. The IR data show a clear dichotomy in H2O detection rate between disks around the lower-mass T Tauri stars and higher-mass, hotter A-type stars
并且因此至少在其表层不干燥。红外数据显示了低质量 T Tauri 恒星周围盘与高质量、热 A 型恒星之间 H 2 O 检测率的明显二分法
(Pontoppidan et al., 2010a; Fedele et al., 2011)
(Pontoppidan 等人,2010a 年; Fedele 等人,2011 年)
. Also, transition disks with inner dust holes show a lack of water line emission. This is likely due to more rapid photodissociation by stars with higher
此外,具有内部尘埃孔的过渡盘显示出缺乏水线发射。这可能是由于更快的光解作用,因为更高的恒星辐射更强
TT_{*}, and thus stronger UV radiation, in regions where the molecules are not shielded by dust.
,因此在分子未被尘埃屏蔽的区域中,UV 辐射更强。

Moving to longer wavelengths, Herschel-PACS spectra probe gas at intermediate radii of the disk, out to 100 AU. Far-IR lines from warm water have been detected in a few disks (Rivière-Marichalar et al., 2012; Meeus et al., 2012; Fedele et al., 2012, 2013). As for the inner disk, the abundance ratios derived from these data are highly uncertain. Sources in which both H2O and CO far-infrared lines have been detected (only a few) indicate H2O/CO column density ratios of 10110^{-1}, suggesting a water abundance of order 10510^{-5} at intermediate layers, but upper limits in other disks suggest values that may be significantly less. Again the disks around T Tauri stars appear to be richer in water than those around A-type stars (Fig. 6).
转移到更长波长,赫歇尔-PACS 光谱探测盘的中间半径处的气体,延伸到 100 天文单位。已在一些盘中检测到来自温暖水的远红外线。至于内盘,从这些数据推导出的丰度比高度不确定。已检测到 H 2 O 和 CO 远红外线的源(仅有少数)表明 H 2 O/CO 柱密度比为 101superscript10110^{-1} ,表明中间层的水丰度约为 105superscript10510^{-5} ,但其他盘的上限值表明可能明显较少。再次,围绕 T Tauri 星的盘似乎比围绕 A 型星的盘富含水(图 6)。

Refer to caption
Figure 6: Near-IR (left) and far-IR (right) spectra of a T Tau and a Herbig Ae disk, showing OH lines in both but H2O primarily in disks around cooler T Tau stars. Figure by D. Fedele, based on Fedele et al. (2011, 2013).
图 6:T Tau 和 Herbig Ae 盘的近红外(左)和远红外(右)光谱,显示两者中的 OH 线,但 H2O 主要存在于较冷的 T Tau 恒星周围的盘中。由 D. Fedele 绘制,基于 Fedele 等人(2011,2013)。

In principle, the pattern of water lines with wavelength should allow the transition from the gaseous water-rich to the water-poor (the snow line) to be probed. As shown by LTE excitation disk models, the largest sensitivity to the location of the snow line is provided by lines in the 40–60 μ\mum region, which is exactly the wavelength range without observational facilities except for SOFIA (Meijerink et al., 2009). For one disk, that around TW Hya, the available shorter and longer wavelength water data have been used to put together a water abundance profile across the entire disk (Zhang et al., 2013). This disk has a dust hole within 4 AU, within which water is found to be depleted. The water abundance rises sharply to a high abundance at the inner edge of the outer disk at 4 AU, but then drops again to very low values as water freezes out in the cold outer disk.
原则上,水线的波长模式应该允许从富含水的气态到贫水(雪线)的过渡得以探测。正如 LTE 激发盘模型所示,对雪线位置的最大敏感性由 40-60 米区域的线提供,这正是除了 SOFIA(Meijerink 等人,2009)之外没有观测设施的波长范围。对于一个盘,即围绕 TW Hya 的盘,已经利用可用的较短和较长波长的水数据来组合整个盘上的水丰度剖面(Zhang 等人,2013)。这个盘在 4 AU 内有一个尘域,其中发现水被耗尽。水丰度在外盘的内边缘(4 AU 处)急剧上升至高丰度,但随后在寒冷的外盘中水结冰后再次降至极低值。

The cold gaseous water reservoir beyond 100 AU is uniquely probed by Herschel-HIFI data of the ground rotational transitions at 557 and 1113 GHz. Weak, but clear detections of both lines have been obtained in two disks, around the nearby T Tau star TW Hya (Hogerheijde et al., 2011) and the Herbig Ae star HD 100546 (Hogerheijde et al., in prep.) (Fig. 3). These are the deepest integrations obtained with the HIFI instrument, with integration times up to 25 hr per line. Similarly deep integrations on 5 other disks do not show detections of water at the same level, nor do shallower observations of a dozen other disks of different characteristics. One possible exception, DG Tau (Podio et al., 2013), is a late class I source with a well-known jet and a high X-ray flux. The TW Hya detection implies abundances of gaseous water around 10710^{-7} in the intermediate layer of the disk, with the bulk of the oxygen in ice on grains at lower layers. Quantitatively, 0.005 Earth oceans of gaseous water and a few thousand oceans of water ice have been detected (1 Earth ocean = 1.4×1024\times 10^{24} gr=0.00023 MEarth). While this is plenty of water to seed an Earth-like planet with water, a single Jovian-type planet formed in this ice-rich region could lock up the bulk of this water.
在距离 100 天文单位之外的冷气态水库被 Herschel-HIFI 数据独特地探测到,这些数据涉及 557 和 1113 GHz 的地面旋转跃迁。在两个盘中分别获得了这两条线的微弱但清晰的探测结果,一个盘围绕附近的 T Tau 星 TW Hya(Hogerheijde 等人,2011)和 Herbig Ae 星 HD 100546(Hogerheijde 等人,准备中)(图 3)。这些是使用 HIFI 仪器获得的最深度的积分,每条线的积分时间长达 25 小时。对其他 5 个盘进行的同样深度的积分未显示出相同水平的水探测结果,对另外十几个具有不同特征的盘的较浅观测也未显示出水的探测结果。一个可能的例外是 DG Tau(Podio 等人,2013),它是一个具有众所周知的喷流和高 X 射线通量的晚 I 类源。TW Hya 的探测结果意味着在盘的中间层周围存在气态水,氧的大部分存在于较低层的颗粒冰中。定量上,已经探测到了 0.005 个地球海洋的气态水和几千个海洋的水冰(1 个地球海洋=1.4 个克=0.00023 M)。 尽管这足以在类地行星上播种水,但在这个富含冰的区域形成的单个类木星型行星可能会锁住大部分水。

Direct detections of water ice are complicated by the fact that IR absorption spectroscopy requires a background light source, and thus a favorable near edge-on orientation of the disk. In addition, care has to be taken that foreground clouds do not contribute to the water ice absorption (Pontoppidan et al., 2005). The 3 μ\mum water ice band has been detected in only a few disks (Terada et al., 2007; Honda et al., 2009). To measure the bulk of the ice, one needs to go to longer wavelengths, where the ice features can be seen in emission. Indeed, the crystalline H2O features at 45 or 60 μ\mum have been detected in several sources with ISO-LWS (Malfait et al., 1998, 1999; Chiang et al., 2001) and Herschel-PACS (McClure et al., 2012, Bouwman et al., in prep). Quantitatively, the data are consistent with 25–50% of the oxygen in water ice on grains in the emitting layer.
水冰的直接探测受到了许多复杂因素的影响,因为红外吸收光谱需要一个背景光源,因此需要一个有利的近边缘方向的盘面。此外,必须小心确保前景云不会对水冰吸收产生影响(Pontoppidan 等人,2005 年)。仅在少数几个盘面中检测到了 3 μ\mu m 水冰带(Terada 等人,2007 年;Honda 等人,2009 年)。要测量冰的大部分,需要转向更长波长,这样冰的特征可以在发射中看到。事实上,ISO-LWS(Malfait 等人,1998 年,1999 年;Chiang 等人,2001 年)和 Herschel-PACS(McClure 等人,2012 年,Bouwman 等人,准备中)已在几个源中检测到了 45 或 60 μ\mu m 处的结晶 H 2 O 特征。定量上,数据与在发射层中颗粒上的水冰中的氧的 25-50%一致。

Refer to caption
Figure 7: Cartoon illustrating the snow line as a function of radius and height in a disk and transport of icy planetesimals across the snowline. Diffusion of water vapor from inner to outer disk followed by freeze-out results in pile-up of ice just beyond the snowline (the cold finger effect). Figure by M. Persson, based on Meijerink et al. (2009); Ciesla and Cuzzi (2006).
图 7:漫画插图,说明盘中的半径和高度作为函数的雪线以及跨越雪线的冰行星体的运输。从内盘向外盘的水蒸气扩散,随后结冰导致冰在雪线之外堆积(冷指效应)。图由 M. Persson 制作,基于 Meijerink 等人(2009);Ciesla 和 Cuzzi(2006)。

The ISO-LWS far-infrared spectra also suggested a strong signature of hydrated silicates in at least one target (Malfait et al., 1999). Newer Herschel-PACS data show no sign of such a feature in the same target (Bouwman, priv. comm.). An earlier claim of hydrated silicates at 2.7 μ\mum in diffuse clouds has now also been refuted (Whittet et al., 1998; Whittet, 2010). Moreover, there is no convincing detection of any mid-infrared feature of hydrated silicates in hundreds of Spitzer spectra of T Tauri (e.g., Olofsson et al., 2009; Watson et al., 2009), Herbig Ae (e.g., Juhász et al., 2010) and warm debris (e.g., Olofsson et al., 2012) disks. Overall, the strong observational consensus is that the silicates prior to planet formation are ‘dry’.
ISO-LWS 远红外光谱还表明至少一个目标中存在水合硅酸盐的强烈特征(Malfait 等人,1999 年)。新的 Herschel-PACS 数据显示在同一目标中没有这种特征(Bouwman,私人通讯)。早期声称弥散云中存在 2.7μm 水合硅酸盐的说法现已被驳斥(Whittet 等人,1998 年;Whittet,2010 年)。此外,在数百个 T Tauri(例如,Olofsson 等人,2009 年;Watson 等人,2009 年)、Herbig Ae(例如,Juhász 等人,2010 年)和温暖碎片(例如,Olofsson 等人,2012 年)盘的 Spitzer 光谱中没有任何水合硅酸盐的中红外特征的令人信服的检测。总体而言,观测结果一致表明在行星形成之前的硅酸盐是“干燥的”。

5.2 Chemical models of disks
5.2 盘的化学模型

The observations of gaseous water discussed in § 5.1 indicate the presence of both rotationally hot Tex450T_{\rm ex}\approx 450 K and cold (Tex<50T_{\rm ex}<50 K) water vapor, with abundances of 104\sim 10^{-4} and much lower values, respectively. Based on the chemistry of water vapor discussed in § 2.4, we expect it to have a relatively well understood distribution within the framework of the disk thermal structure, potentially modified by motions of the various solid or gaseous reservoirs. This is broadly consistent with the observations.
在第 5.1 节讨论的气态水的观测表明存在旋转热 Tex450subscript450T_{\rm ex}\approx 450 K 和冷( Tex<50subscript50T_{\rm ex}<50 K)水蒸气,丰度分别为 104similar-toabsentsuperscript104\sim 10^{-4} 和较低的值。根据第 2.4 节讨论的水蒸气化学,我们预计它在盘状热结构框架内具有相对清晰的分布,可能会受到各种固体或气态储库的运动的影响。这与观测结果基本一致。

Traditionally, the snow line plays a critical role in the distribution of water, representing the condensation or sublimation front of water in the disk, where the gas temperatures and pressures allow water to transition between the solid and gaseous states (Fig. 7). For the solar nebula disk, there is a rich literature on the topic (Hayashi, 1981; Sasselov and Lecar, 2000; Podolak and Zucker, 2004; Lecar et al., 2006; Davis, 2007; Dodson-Robinson et al., 2009). Within our modern astrophysical understanding, this dividing line in the midplane is altered when viewed within the framework of the entire disk physical structure. There are a number of recent models of the water distribution that elucidate these key issues (Glassgold et al., 2009; Woitke et al., 2009b; Bethell and Bergin, 2009; Willacy and Woods, 2009; Gorti et al., 2011; Najita et al., 2011; Vasyunin et al., 2011; Fogel et al., 2011; Walsh et al., 2012; Kamp et al., 2013).
传统上,雪线在水的分布中起着关键作用,代表盘中水的冷凝或升华前缘,在那里气体温度和压力允许水在固体和气态之间转变(图 7)。对于太阳星云盘,关于这个主题有丰富的文献(Hayashi, 1981; Sasselov and Lecar, 2000; Podolak and Zucker, 2004; Lecar et al., 2006; Davis, 2007; Dodson-Robinson et al., 2009)。在我们现代天体物理学的理解中,这个分界线在中面盘内被改变,当在整个盘物理结构的框架内观察时。有许多最近的水分布模型阐明了这些关键问题(Glassgold et al., 2009; Woitke et al., 2009b; Bethell and Bergin, 2009; Willacy and Woods, 2009; Gorti et al., 2011; Najita et al., 2011; Vasyunin et al., 2011; Fogel et al., 2011; Walsh et al., 2012; Kamp et al., 2013)。

5.2.1 General distribution of gaseous water
5.2.1 气态水的一般分布

Fig. 8 shows the distribution of water vapor in a typical kinetic chemical disk model with radius RR and height zz. The disk gas temperature distribution is crucial for the chemistry. It is commonly recognized that dust on the disk surface is warmer than in the midplane due to direct stellar photon heating (Calvet et al., 1992; Chiang and Goldreich, 1997). Furthermore the gas temperature is decoupled from the dust in the upper layers due to direct gas heating (e.g., Kamp and Dullemond, 2004). There are roughly 3 areas where water vapor is predicted to be abundant and therefore possibly emissive (see also discussion in Woitke et al., 2009a). These 3 areas or “regions” are labelled with coordinates (radial and vertical) that are specific to the physical structure (radiation field, temperature, density, dust properties) of this model. Different models (with similar dust- and gas-rich conditions) find the same general structure, but not at the exact same physical location.
图 8 显示了具有半径 RR 和高度 zz 的典型动力学化学盘模型中水蒸气的分布。盘的气体温度分布对化学反应至关重要。众所周知,盘表面的尘埃比中面更暖,这是由于直接的恒星光子加热(Calvet 等,1992; Chiang 和 Goldreich,1997)。此外,由于直接的气体加热(例如,Kamp 和 Dullemond,2004),气体温度与上层尘埃脱耦。大致有 3 个区域被预测为水蒸气丰富且可能具有发射性(另请参阅 Woitke 等人的讨论,2009a)。这 3 个区域或“区域”用特定于该模型的物理结构(辐射场、温度、密度、尘埃特性)的坐标(径向和垂直)标记。不同模型(具有类似的尘埃和气体丰富条件)发现相同的一般结构,但不在完全相同的物理位置。

Region 1 (RR = inner radius to 1.5 AU; z/Rz/R << 0.1): this region coincides with the condensation/sublimation front in the midplane at the snow line. Inside the snow line water vapor will be abundant. Reaction timescales imposed by chemical kinetics limit the overall abundance depending on the gas temperature. As seen in Fig. 8, if the gas temperature exceeds 400\sim 400 K then the midplane water will be quite abundant, carrying all available oxygen not locked in CO and refractory grains. If the gas temperature is below this value, but above the sublimation temperature of 160\sim 160 K, then chemical kinetics could redistribute the oxygen towards other species. During the early gas- and dust-rich stages up to a few Myr, this water vapor dominated region will persist and is seen in nearly all models. However, as solids grow, the penetrating power of UV radiation is increased. Since water vapor is sensitive to photodissociation by far-UV, this could lead to gradual decay of this layer, which would be consistent with the non-detection of water vapor inside the gaps of a small sample of transition disks (e.g., Pontoppidan et al., 2010a; Zhang et al., 2013).
区域 1( RR = 内半径至 1.5 AU; z/Rz/R << 0.1):该区域与中面板上的凝结/升华前沿在雪线处重合。在雪线内部,水蒸气将会很丰富。由化学动力学所施加的反应时间尺度限制了整体丰度,取决于气体温度。如图 8 所示,如果气体温度超过 400similar-toabsent400\sim 400 K,则中面板水将会相当丰富,携带所有未锁定在 CO 和耐火颗粒中的氧。如果气体温度低于此值,但高于 160similar-toabsent160\sim 160 K 的升华温度,则化学动力学可能会将氧重新分配到其他物种。在早期气体和富含尘埃的阶段,长达几百万年,这个水蒸气主导的区域将持续存在,并且几乎在所有模型中都能看到。然而,随着固体的增长,紫外辐射的穿透能力增强。由于水蒸气对远紫外的光解作用敏感,这可能导致该层逐渐衰减,这与一小部分过渡盘中的水蒸气未被检测到是一致的(例如 Pontoppidan 等人,2010a;Zhang 等人,2013)。

Region 2 (R>R> 20 AU; surface layers and outer disk midplane): In these disk layers the dust temperature is uniformly below the sublimation temperature of water. Furthermore at these high densities (n>106n>10^{6} cm-3) atoms and molecules freeze out on dust grains on short timescales (§2.4). Under these circumstances, in the absence of non-thermal desorption mechanisms, models predict strong freeze-out with the majority of available oxygen present on grains as water ice. Much of this may be primordial water ice supplied by the natal cloud (Visser et al., 2011, Fig. 5).
区域 2( R>absentR> 20 AU;表面层和外部盘面):在这些盘层中,尘埃温度均低于水的升华温度。此外,在这些高密度( n>106superscript106n>10^{6} cm -3 )的情况下,原子和分子会在短时间内冻结在尘埃颗粒上(§2.4)。在这种情况下,在非热解吸机制的缺乏下,模型预测会出现强烈的冻结,大部分氧会以水冰的形式存在于颗粒上。其中很多可能是出生云提供的原始水冰(Visser 等人,2011 年,图 5)。

The detection of rotationally cold water vapor emission in the outer disk of TW Hya demonstrates that a tenuous layer of water vapor is present and that some non-thermal desorption process is active (Hogerheijde et al., 2011). The leading candidate is photodesorption of water ice (Dominik et al., 2005; Öberg et al., 2009), as discussed in § 2.4.4, particularly given the high UV luminosities of T Tauri stars (Yang et al., 2012a). This UV excess is generated by accretion and dominated by Lyα\alpha line emission (Schindhelm et al., 2012).

Once desorbed as OH and H2O, the UV radiation then also destroys the water vapor molecules leading to a balance between these processes and a peak abundance near (13)×107(1-3)\times 10^{-7} (Dominik et al., 2005; Hollenbach et al., 2009). In general most models exhibit this layer, which is strongly dependent on the location and surface area of ice-coated grains (i.e. less surface area reduces the effectiveness of photodesorption). Direct comparison of models with observations finds that the amount of water vapor predicted to be present exceeds the observed emission (Bergin et al., 2010; Hogerheijde et al., 2011). This led to the suggestion that the process of grain growth and sedimentation could operate to remove water ice from the UV exposed disk surface layers. This is consistent with spectroscopic data of the TW Hya scattered light disk, which do not show water ice features in the spectrum originating from this layer (Debes et al., 2013). However, further fine tuning of this settling mechanism is needed (Dominik and Dullemond, in prep., Akimkin et al., 2013). An alternative explanation may be a smaller dust disk compared with the gas disk (Qi et al., 2013).
一旦作为 OH 和 H 2 O 脱附,紫外辐射也会破坏水蒸气分子,导致这些过程之间的平衡,并在 (13)×10713superscript107(1-3)\times 10^{-7} 附近达到峰值丰度(Dominik 等,2005 年;Hollenbach 等,2009 年)。一般来说,大多数模型都展示了这一层,这一层强烈依赖于冰覆盖颗粒的位置和表面积(即表面积较小会降低光解吸效率)。模型与观测数据的直接比较发现,预测存在的水蒸气量超过了观测到的发射量(Bergin 等,2010 年;Hogerheijde 等,2011 年)。这导致了一种观点,即颗粒生长和沉积的过程可能会移除受紫外线暴露的盘面层中的水冰。这与 TW Hya 散射光盘的光谱数据一致,该光谱数据不显示来自该层的水冰特征(Debes 等,2013 年)。然而,需要进一步对这种沉降机制进行微调(Dominik 和 Dullemond,准备中,Akimkin 等,2013 年)。另一种解释可能是尘埃盘较气体盘更小(Qi 等,2013 年)。

Refer to caption
Figure 8: Abundance of gaseous water relative to total hydrogen as a function of radial distance, RR, and relative height above the midplane, z/Rz/R, for a disk around an A-type star (TT_{*}=8600 K). Three regions with high H2O abundance can be distinguished. Regions 1 and 3 involve high-temperature chemistry, whereas region 2 lies beyond the snow line and involves photodesorption of water ice. The white contours indicate gas temperatures of 200 and 1500 K, whereas the red contour shows the nH=5×1010n_{\rm H}=5\times 10^{10} cm-3 density contour. From Woitke et al. (2009b).
图 8:在 A 型恒星周围的盘中,相对于总氢的径向距离 RR 和相对于中平面的高度 z/Rz/R 的气态水丰度。可以区分出三个具有高 H 2 O 丰度的区域。区域 1 和 3 涉及高温化学反应,而区域 2 位于雪线之外,涉及水冰的光解吸附。白色轮廓表示气体温度为 200 和 1500 K,而红色轮廓显示 nH=5×1010subscript5superscript1010n_{\rm H}=5\times 10^{10} cm -3 密度轮廓。来自 Woitke 等人(2009b)。

Region 3 (R<R< 20 AU; z/R>z/R> 0.1): Closer to the exposed disk surface the gas and dust become thermally decoupled. The density where this occurs depends on the relative amount of dust grains in the upper atmosphere, which may be altered by dust coagulation and settling (Jonkheid et al., 2004; Nomura et al., 2007) and on the thermal accommodation of the dust gas interaction (Burke and Hollenbach, 1983). In these decoupled layers TgasTdustT_{\rm gas}\gg T_{\rm dust}, and when the gas temperature exceeds a few hundred K the neutral-neutral gas-phase pathways for water formation become efficient, leading to water abundances of order 10510^{-5} (Fig. 8).
区域 3( R<absentR< 20 AU; z/R>absentz/R> 0.1):靠近暴露的盘面表面,气体和尘埃变得热力解耦。这种情况发生的密度取决于上层大气中尘埃颗粒的相对数量,这可能会受到尘埃凝聚和沉降(Jonkheid 等人,2004 年;Nomura 等人,2007 年)以及尘埃气体相互作用的热适应(Burke 和 Hollenbach,1983 年)的影响。在这些解耦层中 TgasTdustmuch-greater-thansubscriptsubscriptT_{\rm gas}\gg T_{\rm dust} ,当气体温度超过几百 K 时,水形成的中性-中性气相途径变得高效,导致水的丰度约为 105superscript10510^{-5} (图 8)。

More directly, the disk surface is predicted to be water vapor rich at gas temperatures \gtrsim few hundred K and dust temperatures \sim100 K. Indeed, there should exist surface layers at radii where the midplane temperature is sufficiently low to freeze water vapor, but where the surface can support water formation via the high-temperature chemistry (i.e., region 3 goes out to larger radii than region 1). Thus the water zone on the disk surface presents the largest surface area and it is this water that is readily detected with current astronomical observations of high-lying transitions of H162{}_{2}^{16}O with Spitzer and Herschel. The snow line in the midplane is thus potentially hidden by the forest of water transitions produced by the hot chemistry on the surface.
更直接地,盘面被预测为在气体温度约为几百 K 和尘埃温度约为 100K 时富含水蒸气。事实上,在中面温度足够低以冻结水蒸气的半径处应存在表面层,但表面可以通过高温化学来支持水的形成(即,区域 3 的半径比区域 1 大)。因此,盘面上的水区域呈现出最大的表面积,而正是这种水可以通过目前的高位 H2O 的 Spitzer 和 Herschel 的天文观测轻松检测到。中面的雪线因此可能被表面上的热化学产生的水转变的森林所隐藏。

Refer to caption
Figure 9: Cartoon illustrating the water self-shielding mechanism and the resulting vertical stratification of O and H2O. Inclusion of water self-shielding in the upper layers leads to a ‘wet’ warm layer. From Bethell and Bergin (2009).
图 9:漫画描绘了水自屏蔽机制以及 O 和 H2O 的垂直分层。在上层包含水自屏蔽会导致一个“湿润”的温暖层。来自 Bethell 和 Bergin(2009)。

There are some key dependences and differences which can be highlighted. One important factor is the shape of the UV radiation field. In general, models that use a scaled interstellar UV radiation field, for example based on FUSE/IUE/HST observations of the UV excess (Yang et al., 2012a), neglect the fact that some molecules like H2 and CO require very energetic photons to photodissociate, which are not provided by very cool stars. A better approach is to take the actual stellar continua into account (van Dishoeck et al., 2006), with UV excess due to accretion added where appropriate (van Zadelhoff et al., 2003). A very important factor in this regard is the relative strength of the Lyα\alpha line to the overall UV continuum. Observations find that Lyα\alpha has nearly an order of magnitude more UV flux than the stellar FUV continuum in accreting sources (Bergin et al., 2003; Schindhelm et al., 2012). In addition because of the difference in scattering (Lyα\alpha isotropic from H atom surface; UV continuum anisotropic from dust grains), Lyα\alpha will dominate the radiation field deeper into the disk (Bethell and Bergin, 2011).
一些关键依赖关系和差异可以被突出显示。一个重要因素是紫外辐射场的形状。通常,使用缩放的星际紫外辐射场的模型,例如基于 FUSE/IUE/HST 对紫外线过剩的观测(Yang 等,2012a),忽略了一些分子如 H 2 和 CO 需要非常高能的光子来光解,这些光子不是由非常冷的恒星提供的。更好的方法是考虑实际的恒星连续体(van Dishoeck 等,2006),在适当的地方添加由吸积引起的紫外线过剩(van Zadelhoff 等,2003)。在这方面一个非常重要的因素是 Ly α\alpha 线相对于整体紫外线连续体的强度。观测发现,在吸积源中,Ly α\alpha 的紫外线通量几乎比恒星 FUV 连续体高一个数量级(Bergin 等,2003;Schindhelm 等,2012)。此外,由于散射的差异(Ly α\alpha 来自 H 原子表面的各向同性;紫外线连续体来自尘埃颗粒的各向异性),Ly α\alpha 将在盘面深处主导辐射场(Bethell 和 Bergin,2011)。

Most models find the presence of this warm water layer in dust-dominated disks. However, Bethell and Bergin (2009) suggest that water can form in such high abundances in the surface layer that it mediates the transport of the energetic UV radiation by becoming self-shielding (Fig. 9). If this is the case, then the surface water would survive for longer timescales, because it is somewhat decoupled from the dust evolution. As a consequence water and chemistry in the midplane might be protected even as the FUV absorbing dust grains settle to the midplane.
大多数模型发现尘埃主导的盘中存在这一温暖水层。然而,Bethell 和 Bergin(2009)提出,水可以在表层形成如此丰富,以至于通过自我屏蔽介导能量紫外辐射的传输(图 9)。如果是这种情况,那么表层水将存活更长时间,因为它在一定程度上与尘埃演化脱钩。因此,即使 FUV 吸收尘埃颗粒沉降到中平面,中平面中的水和化学物质可能仍受到保护。

Additional factors of importance for the survival of this emissive surface layer are the gas temperature and molecular hydrogen abundance. Although theoretical solutions for the gas temperature are inherently uncertain, it is clear that hot (Tgas>T_{\rm gas}> few hundred K) layers exist on disk surfaces (Bruderer et al., 2012). However, as the gas disk dissipates, the accretion rate onto the star decays on timescales of a few Myr. Thus the UV luminosity that is associated with this accretion declines and the disk will cool down, cutting off the production of water from the hot (TT\gtrsim 400 K) gas-phase chemistry on the exposed surface. In addition, as shown by Glassgold et al. (2009) and Ádámkovics et al. (2013), the formation of surface water requires the presence of H2 to power the initiating reaction. Finally, vertical mixing through turbulence or disk winds can bring water ice from the lower to the upper layers where the ice sublimates and adds to the oxygen budget and water emission (Heinzeller et al., 2011).
重要因素之一是这种发射表面层的存活的气体温度和分子氢丰度。尽管气体温度的理论解决方案本质上是不确定的,但很明显,在盘面上存在着热( Tgas>subscriptabsentT_{\rm gas}> 几百 K)的层(Bruderer 等,2012)。然而,随着气体盘逐渐消散,星体上的吸积速率会在几百万年的时间尺度上衰减。因此,与这种吸积相关的紫外线亮度会下降,盘会冷却下来,切断了来自暴露表面上热( Tgreater-than-or-equivalent-toabsentT\gtrsim 400 K)气相化学的水的产生。此外,正如 Glassgold 等人(2009 年)和Ádámkovics 等人(2013 年)所示,表面水的形成需要 H 2 的存在来驱动起始反应。最后,通过湍流或盘风的垂直混合可以将水冰从下层带到上层,水冰升华并增加氧的预算和水的排放(Heinzeller 等,2011)。

5.2.2 Planetesimal formation and water ice transport
5.2.2 行星体形成和水冰运输

The dust particles in disks collide and grow, with water ice mantles generally thought to help the coagulation processes. The evolution of dust to pebbles, rocks and planetesimals (1–100 km bodies, the precursors of comets and asteroids) is described in the chapters by Testi et al. and Johansen et al.. The disk models cited above do not take into account transport of ice-rich planetesimals from the cold outer to the warm inner disk, even though such radial drift is known to be highly effective at a few AU for rocks up to meter size (or mm size further out in the disk) (e.g., Weidenschilling and Cuzzi, 1993). This drift of icy planetesimals can be a source of water vapor enrichment inside the snow line (Ciesla and Cuzzi, 2006). The astrophysical signature of this phenomenon would be the presence of water vapor in the inner disk with an abundance greater than the stellar oxygen abundance, because the planetesimals are hydrogen-poor, that is, the main volatile species, H2, is not present in water-rich planetesimal ices. Some of this hot water can diffuse outwards again and re-condense just outside the snowline (the cold finger effect. Fig. 7) which increases the density of solids by a factor of 2–4 and thereby assists planet formation (Stevenson and Lunine, 1988). Alternatively, icy grains can be trapped in pressure bumps where they can grow rapidly to planetesimals before moving inward (e.g., van der Marel et al., 2013).
盘中的尘埃颗粒相互碰撞并增长,通常认为水冰外壳有助于凝聚过程。从尘埃到鹅卵石、岩石和小行星(1-100 公里大小的天体,是彗星和小行星的前体)的演化在 Testi 等人和 Johansen 等人的章节中有所描述。上述引用的盘模型并未考虑将富含冰的小行星从寒冷的外部运输到温暖的内部盘中,尽管已知这种径向漂移对于直径达米级(或盘外更远处的毫米级)的岩石在几个天文单位内是非常有效的(例如,Weidenschilling 和 Cuzzi,1993 年)。这些冰冷小行星的漂移可能是内部雪线内水蒸气富集的来源(Ciesla 和 Cuzzi,2006 年)。这种现象的天体物理特征将是内部盘中存在水蒸气,其丰度大于恒星氧的丰度,因为小行星缺乏氢,也就是说,主要的挥发性物种 H 2 不会存在于富含水的小行星冰中。其中一些热水可能再次向外扩散并在雪线外重新凝结(冷指效应。图。 7)这会将固体的密度增加 2-4 倍,从而促进行星形成(Stevenson 和 Lunine,1988 年)。或者,冰颗粒可以被困在压力隆起中,在那里它们可以在向内移动之前迅速增长为小行星(例如,van der Marel 等人,2013 年)。

6 WATER IN THE OUTER SOLAR SYSTEM
6 外太阳系中的水

In the standard model of the disk out of which our solar system formed, the snow line was at 2.7 AU at the end of the gas-rich phase (Hayashi, 1981). This snow line likely moved inward from larger distances in the early embedded phase (Kennedy and Kenyon, 2008). Thus, it is no surprise that water ice is a major constituent of all solar system bodies that formed and stayed beyond the snow line. Nevertheless, their water ice content, as measured by the mass in ice with respect to total ice+rock mass can differ substantially, from 1% for some asteroids to typically 50% for comets. Also, the observational signatures of water on these icy bodies and its isotopic ratio can differ. In the following sections, we review our knowledge of water ice in the present-day solar system. In § 8, possible mechanisms of supplying water from these reservoirs to the terrestrial planet zone will be discussed.<<
在我们太阳系形成的盘状标准模型中,雪线在富含气体的末期位于 2.7 天文单位处(Hayashi, 1981)。这条雪线可能在早期嵌入阶段从更远的距离向内移动(Kennedy 和 Kenyon, 2008)。因此,毫不奇怪,水冰是所有形成并留在雪线之外的太阳系天体的主要组成部分。然而,它们的水冰含量,即以冰质量占总冰+岩石质量的比例,可以有很大差异,从一些小行星的 1%到彗星通常的 50%。此外,这些冰天体上水的观测特征及其同位素比率也可能不同。在接下来的章节中,我们将回顾我们对现今太阳系中水冰的了解。在第 8 节中,将讨论从这些储库向地球型行星区域供应水的可能机制。 <<

6.1 Outer asteroid belt  6.1 外小行星带

Since the outer asteroid belt is located outside the Hayashi snow line, it provides a natural reservoir of icy bodies in the solar system. This part of the belt is dominated by so-called C, P and D class asteroids with sizes up to a few 100 km at distances of \sim3, 4 and \geq4 AU, respectivily, characterized by their particularly red colors and very low albedos, <{{}_{<}\atop{{}^{\sim}}}0.1 (Bus and Binzel, 2002). Because of their spectroscopic similarities to the chemically primitive carbonaceous chondrites found as meteorites on Earth, C-type asteroids have been regarded as largely unaltered, volatile-rich bodies. The P- and D-types may be even richer in organics.
由于外部小行星带位于哈亚希雪线之外,它为太阳系提供了一个天然的冰体储库。这部分小行星带主要由所谓的 C、P 和 D 类小行星主导,其大小可达几百公里,在距离分别为 3、4 和 4 天文单位的位置上,其特点是特别红色和非常低的反照率,为 0.1(Bus and Binzel, 2002)。由于它们在光谱上与地球上发现的化学原始的含碳球粒陨石相似,C 型小行星被认为是基本未改变的、富含挥发性物质的天体。P 型和 D 型可能甚至更富含有机物质。

The water content in these objects has been studied though IR spectroscopy of the 3 μ\mum band. In today’s solar system, any water ice on the surface would rapidly sublimate at the distance of the belt, so only water bonded to the rocky silicate surface is expected to be detected. Hydrated minerals can be formed if the material has been in contact with liquid water. The majority of the C-type asteroids show hydrated silicate absorption, indicating that they indeed underwent heating and aqueous alteration episodes (Jones et al., 1990, and refs. therein). However, only 10% of the P and D-type spectra show weak water absorption, suggesting that they have largely escaped this processing and that the abundance of hydrated silicates gradually declines in the outer asteroid belt. Nevertheless, asteroids that do not display water absorption on their surfaces (mostly located beyond 3.5 AU) may still retain ices in their interior. Indeed, water fractions of 5–10% of their total mass have been estimated. This is consistent with models that show that buried ice can persist in the asteroid belt within the top few meters of the surface over billions of years, as long as the mean surface temperature is less than about 145 K (Schorghofer, 2008). Water vapor has recently been detected around the dwarf planet Ceres at 2.7 AU in the asteroid belt, with a production rate of at least 102610^{26} mol s-1, directly confirming the presence of water (Küppers et al., 2014).
这些天体中的水含量已通过 3 μ\mu m 带的红外光谱研究。在今天的太阳系中,表面上的任何水冰在带的距离处会迅速升华,因此预计只有与岩石硅酸盐表面结合的水才能被检测到。如果物质接触过液态水,则可以形成含水矿物。大多数 C 型小行星显示出含水硅酸盐吸收,表明它们确实经历了加热和水蚀变事件(Jones 等人,1990 年,以及其中的参考文献)。然而,只有 10%的 P 型和 D 型光谱显示出弱的水吸收,这表明它们在很大程度上逃脱了这种处理,而含水硅酸盐的丰度在外部小行星带中逐渐下降。然而,表面上不显示水吸收的小行星(主要位于 3.5 AU 之外)可能仍然在其内部保留着冰。事实上,它们的总质量中已估计出 5-10%的水分。 这与模型一致,该模型显示,在表面顶部的几米深处,只要平均表面温度低于约 145K(Schorghofer,2008 年),在数十亿年内,埋藏的冰可以在小行星带中持续存在。最近在小行星带 2.7 天文单位处探测到了水蒸气,其产生速率至少为 1026superscript102610^{26} 摩尔每秒 -1 ,直接证实了水的存在(Küppers 等,2014 年)。

Hsieh and Jewitt (2006) discovered a new population of small objects in the main asteroid belt, displaying cometary characteristics. These so-called main belt comets (ten are currently known) display clearly asteroidal orbits, yet have been observed to eject dust and thus satisfy the observational definition of a comet. These objects are unlikely to have originated elsewhere in the solar system and to have subsequently been trapped in their current orbits. Instead, they are intrinsically icy bodies, formed and stored at their current locations, with their cometary activity triggered by some recent event.
Hsieh 和 Jewitt(2006 年)发现了主小行星带中一群新的小天体,显示出彗星特征。这些所谓的主带彗星(目前已知十个)明显显示出小行星的轨道,但已观察到它们喷射尘埃,因此满足了彗星的观测定义。这些天体不太可能起源于太阳系中的其他地方,然后被困在它们当前的轨道中。相反,它们是固有的冰体,在它们当前的位置形成并储存,其彗星活动是由某些最近的事件触发的。

Since main belt comets are optically faint, it is not known whether they display hydration spectral features that could point to the presence of water. Activity of main belt comets is limited to the release of dust and direct outgassing of volatiles, like for Ceres, has so far not been detected. The most stringent indirect upper limit for the water production rate derived from CN observations is that in the prototypical main belt comet 133P/Elst-Pizarro <1.3×1024<1.3\times 10^{24} mol s-1 (Licandro et al., 2011), which is subject to uncertainties in the assumed water-to-CN abundance ratio. For comparison, this is five orders of magnitude lower than the water production rate of comet Hale-Bopp. Its mean density is 1.3 gr cm-3 suggesting a moderately high ice fraction (Hsieh et al., 2004). Herschel provided the most stringent direct upper limits for water outgassing in 176P/LINEAR (<4×1025<4\times 10^{25} mol s-1, 3σ\sigma; de Val-Borro et al. 2012) and P/2012 T1 PANSTARRS (<8×1025<8\times 10^{25} mol s-1; O’Rourke et al. 2013).
由于主带彗星在光学上很暗淡,目前尚不清楚它们是否显示出可能指向水存在的水合物光谱特征。主带彗星的活动仅限于释放尘埃和挥发物的直接放气,就像矮行星谷神星一样,迄今为止尚未检测到。从 CN 观测中推导出的水产生速率的最严格的间接上限是原型主带彗星 133P/Elst-Pizarro 的 <1.3×1024absent1.3superscript1024<1.3\times 10^{24} mol s -1 (Licandro 等,2011 年),这取决于假定的水对 CN 丰度比的不确定性。作为比较,这比哈雷-博普彗星的水产生速率低五个数量级。其平均密度为 1.3 gr cm -3 ,表明具有适度高的冰含量(Hsieh 等,2004 年)。赫歇尔提供了 176P/LINEAR( <4×1025absent4superscript1025<4\times 10^{25} mol s -1 ,3 σ\sigma ;de Val-Borro 等,2012 年)和 P/2012 T1 PANSTARRS( <8×1025absent8superscript1025<8\times 10^{25} mol s -1 ;O’Rourke 等,2013 年)水放气的最严格的直接上限。

Another exciting discovery is the direct spectroscopic detection of water ice on the asteroid 24 Themis (Campins et al., 2010; Rivkin and Emery, 2010), the largest (198 km diameter) member of the Themis dynamical family at \sim3.2 AU, which also includes three main belt comets. The 3.1 μ\mum spectral feature detected in Themis is significantly different from those in other asteroids, meteorites and all plausible mineral samples available. Campins et al. (2010) argue that the observations can be accurately matched by small ice particles evenly distributed on the surface. A subsurface ice reservoir could also be present if Themis underwent differentiation resulting in a rocky core and an ice mantle. Jewitt and Guilbert-Lepoutre (2012) find no direct evidence of outgassing from the surface of Themis or Cybele with a 5σ\sigma upper limit for the water production rate 1.3×10281.3\times 10^{28} mol s-1, assuming a cometary water-to-CN mixing ratio. They conclude that any ice that exists on these bodies should be relatively clean and confined to a 10% fraction of the Earth-facing surface.<<
另一个令人兴奋的发现是直接光谱检测到小行星 24 Themis 上的水冰(Campins 等,2010 年;Rivkin 和 Emery,2010 年),它是 Themis 动力家族中最大的成员(直径 198 公里),位于 3.2 AU 处,该家族还包括三颗主带彗星。在 Themis 上检测到的 3.1 米光谱特征与其他小行星、陨石和所有可行的矿物样品中的特征明显不同。Campins 等人(2010 年)认为,这些观测结果可以通过表面均匀分布的小冰颗粒准确匹配。如果 Themis 经历了分化,形成了岩质核心和冰外壳,那么地下冰库也可能存在。Jewitt 和 Guilbert-Lepoutre(2012 年)未发现来自 Themis 或 Cybele 表面的气体排放的直接证据,假设彗星水与 CN 混合比的上限为 5 个 mol/s,他们得出结论,这些天体上存在的任何冰应该相对干净,并且局限于地球朝向表面的 10%部分。

Altogether, these results suggest that water ice may be common below asteroidal surfaces and widespread in asteroidal interiors down to smaller heliocentric distances than previously expected. Their water contents are clearly much higher than those of meteorites that originate from the inner asteroid belt, which have only 0.01% of their mass in water (Hutson and Ruzicka, 2000).
总的来说,这些结果表明,在小行星表面以下,水冰可能很常见,并且在小行星内部的分布范围比先前预期的更小。它们的含水量明显比源自内小行星带的陨石高得多,后者只有 0.01%的质量是水(Hutson 和 Ruzicka,2000)。

6.2 Comets  彗星

Comets are small solar system bodies with radii less than 20 km that have formed and remained for most of their lifetimes at large heliocentric distances. Therefore, they likely contain some of the least-processed, pristine ices from the solar nebula disk. They have often been described as ‘dirty snowballs’, following the model of Whipple (1950), in which the nucleus is visualized as a conglomerate of ices, such as water, ammonia, methane, carbon dioxide, and carbon monoxide, combined with meteoritic materials. However, Rosetta observations during the Deep Impact encounter (Küppers et al., 2005) suggest a dust-to-gas ratio in excess of unity in comet 9P/Tempel 1. Typically, cometary ice/rock ratios are of order unity with an implied porosity well over 50% (A’Hearn, 2011).
彗星是太阳系中半径小于 20 公里的小天体,它们在大半径处形成并在大部分寿命中保持。因此,它们可能含有太阳星云盘中一些未经加工的原始冰。它们经常被描述为“脏雪球”,遵循 Whipple(1950)的模型,其中核被视为由水、氨、甲烷、二氧化碳和一氧化碳等冰以及陨石材料组成的凝聚物。然而,罗塞塔在深度撞击事件期间的观测(Küppers 等,2005)表明,彗星 9P/Tempel 1 中的尘埃与气体比例超过了单位。通常,彗星的冰/岩比例约为单位数量级,暗示孔隙度远远超过 50%(A’Hearn,2011)。

The presence and amount of water in comets is usually quantified through their water production rates, which are traditionally inferred from radio observations of its photodissociation product, OH, at 18 cm (Crovisier et al., 2002). Measured rates vary from 102610^{26} to 102910^{29} mol s-1. The first direct detection of gaseous water in comet 1P/Halley, through its ν3\nu_{3} vibrational band at 2.65 μ\mum, was obtained using the KAO (Mumma et al., 1986). The 557 GHz transition of ortho-water was observed by SWAS (Neufeld et al., 2000) and OdinOdin (Lecacheux et al., 2003), whereas Herschel provided for the first time access to multiple rotational transitions of both ortho- and para-water (Hartogh et al., 2011). These multi-transition mapping observations show that the derived water production rates are sensitive to the details of the excitation model used, in particular the ill-constrained temperature profile within the coma, with uncertainties up to 50% (Bockelée-Morvan et al., 2012).
彗星中水的存在和数量通常是通过其水产生速率来量化的,这些速率传统上是通过其光解产物 OH 的无线电观测来推断的,观测波长为 18 厘米(Crovisier 等,2002 年)。测量速率从 0 到 1 摩尔每秒不等。通过其 2.65 微米处的 3 个振动带,首次直接探测到 1P/哈雷彗星中的气态水,使用了 KAO(Mumma 等,1986 年)。SWAS(Neufeld 等,2000 年)和 OdinOdin (Lecacheux 等,2003 年)观测到了正水的 557 GHz 跃迁,而赫歇尔首次提供了对正水和顺水的多个旋转跃迁的访问(Hartogh 等,2011 年)。这些多跃迁映射观测表明,得出的水产生速率对所使用的激发模型的细节敏感,特别是对于尾气中温度剖面的不确定性高达 50%(Bockelée-Morvan 等,2012 年)。

Dynamically, comets can be separated into two general groups: short-period, Jupiter-family comets and long-period comets (but see Horner et al. 2003 for a more detailed classification). Short-period comets are thought to originate from the Kuiper belt, or the associated scattered disk, beyond the orbit of Neptune, while long-period comets formed in the Jupiter-Neptune region and were subsequently ejected into the Oort cloud by gravitational interactions with the giant planets. In reality, the picture is significantly more complex due to migration of the giant planets in the early solar system (see below). In addition, recent simulations (Levison et al., 2010) suggest that the Sun may have captured comets from other stars in its birth cluster. In this case, a substantial fraction of the Oort-cloud comets, perhaps in excess of 90%, may not even have formed in the Sun’s protoplanetary disk. Consequently, there is increasing emphasis on classifying comets based on their chemical and isotopic composition rather than orbital dynamics (Mumma and Charnley, 2011). There is even evidence for heterogeneity within a single comet, illustrating that comets may be built up from cometesimals originating at different locations in the disk (A’Hearn, 2011).
动态地,彗星可以分为两个一般性群体:短周期、木星家族彗星和长周期彗星(但请参见 Horner 等人 2003 年的更详细分类)。短周期彗星被认为起源于冥王星轨道外的柯伊伯带或相关的散射盘,而长周期彗星形成于木星-海王星区域,随后通过与巨行星的引力相互作用被抛入奥尔特云。实际上,由于早期太阳系中巨行星的迁移(见下文),情况要复杂得多。此外,最近的模拟(Levison 等人,2010 年)表明太阳可能从其诞生星团中捕获了其他恒星的彗星。在这种情况下,奥尔特云彗星的相当一部分,也许超过 90%,甚至可能并非在太阳的原行星盘中形成。因此,越来越强调根据彗星的化学和同位素组成而不是轨道动力学对彗星进行分类(Mumma 和 Charnley,2011)。 甚至有证据表明单个彗星内部存在异质性,说明彗星可能是由来自盘中不同位置的彗核物质组成的(A'Hearn, 2011)。

Traditionally, the ortho-to-para ratio in water and other cometary volatiles has been used to contrain the formation and thermal history of the ices. Recent laboratory experiments suggest that this ratio is modified by the desorption processes, both thermal sublimation and photodesorption, and may therefore tell astronomers less about the water formation location than previously thought (see discussions in van Dishoeck et al. 2013 and Tielens 2013).
传统上,水和其他彗星挥发物中的正对位比已被用来限制冰的形成和热史。最近的实验室实验证明,这种比率受到脱附过程的影响,包括热升华和光解吸,因此可能告诉天文学家的信息比之前认为的关于水形成位置少(请参见 van Dishoeck 等人 2013 年和 Tielens 2013 年的讨论)。

6.3 Water in the outer satellites
外卫星中的水

Water is a significant or major component of almost all moons of the giant planets for which densities or spectral information are available. Jupiter’s Galilean moons exhibit a strong gradient from the innermost (Io, essentially all rock, no ice detected) to the outermost (Callisto, an equal mixture of rock and ice). Ganymede has nearly the same composition and hence rock-to-ice ratio as Callisto. Assuming that the outermost of the moons reflects the coldest part of the circumplanetary disk out of which the moons formed, and hence full condensation of water, Callisto’s (uncompressed) density matches that of solar-composition material in which the dominant carbon-carrier was methane rather than carbon monoxide (Wong et al., 2008). The fact that Callisto has close to (but not quite) the full complement of water expected based on the solar oxygen abundance implies that the disk around Jupiter had a different chemical composition (CH4-rich, CO-poor) from that of the solar nebula disk, which was CO-dominated (Prinn and Fegley, 1981).
水是几乎所有巨大行星的卫星中的一个重要或主要组成部分,只要有密度或光谱信息可用。 木星的伽利略卫星显示出从最内侧(Io,基本上全是岩石,未检测到冰)到最外侧(木卫一,岩石和冰的混合物)的强烈梯度。 木卫三的成分几乎与木卫一相同,因此岩石与冰的比率也相同。 假设卫星中最外侧反映了卫星形成的环行星盘的最冷部分,因此水完全凝结,木卫一(未压缩)的密度与太阳组成物质的密度相匹配,其中主要的碳载体是甲烷而不是一氧化碳(Wong 等,2008 年)。 木卫一几乎具有预期的太阳氧丰度所期望的全部水含量,这意味着木星周围的盘具有不同的化学组成(富含 CH 4 ,贫含 CO),而不是太阳星云盘的 CO 主导(Prinn 和 Fegley,1981 年)。

The Saturnian satellites are very different. For satellites large enough to be unaffected by porosity, but excluding massive Titan, the ice-to-rock ratio is higher than for the Galilean moons (Johnson and Estrada, 2009). However, Titan—by far the most massive moon—has a bulk density and mass just in between, and closely resembling, Ganymede’s and Callisto’s. Evidently the Saturnian satellite system had a complex collisional history, in which the original ice-rock ratio of the system was not preserved except perhaps in Titan. The moon Enceladus, which exhibits volcanic and geyser activity, offers the unique opportunity to sample Saturnian system water directly. Neptune’s Triton, like Pluto, has a bulk density and hence ice-rock ratio consistent with what is expected for a solar nebula disk in which CO dominated over CH4. Its water fraction is about 15–35%. At this large distance from the Sun, N2 can also be frozen out and Triton’s spectrum is indeed dominated by N2 ice with traces of CH4 and CO ices; the water signatures are much weaker than on other satellites (Cruikshank et al., 2000). The smaller Trans Neptunion Objects (TNOs) (few hundred km size) are usually found to have mean densities around 1 g cm-3 and thus a high ice fraction. Larger TNOs such as Quaoar and Haumea have much higher densities (2.6–3.3 g cm-3) suggesting a much lower ice content, even compared with Pluto (2.0 g cm-3<<) (Fornasier et al., 2013).
土星的卫星非常不同。对于足够大以不受多孔性影响的卫星,但不包括质量巨大的泰坦,其冰岩比高于伽利略卫星(Johnson and Estrada, 2009)。然而,泰坦——迄今为止最大的卫星——具有介于甘尼米德和卡利斯托之间的体积密度和质量,非常类似。显然,土星的卫星系统经历了复杂的碰撞历史,原始冰岩比可能仅在泰坦中得以保留。展示火山和间歇泉活动的卫星土卫二提供了直接采样土星系统水的独特机会。海王星的特里同和冥王星一样,具有符合太阳星云盘中 CO 主导 CH 的体积密度和因此冰岩比的特征。其水分约为 15-35%。在离太阳如此遥远的地方,氮也可能被冻结,特里同的光谱确实以 N2 冰为主,带有少量 CH3 和 CO 冰;水的特征要比其他卫星弱得多。 较小的距海王星天体(TNOs)(几百公里大小)通常具有平均密度约为 1 克/立方厘米,因此含有很高的冰含量。较大的 TNOs,如夸奥尔和哈美亚,具有更高的密度(2.6-3.3 克/立方厘米),表明冰含量较低,甚至与冥王星(2.0 克/立方厘米)相比也是如此(Fornasier 等人,2013 年)。

In summary, the water ice content of the outer satellites varies with position and temperature, not only as a function of distance from the Sun, but also from its parent planet. Ice fractions are generally high (\geq50%) in the colder parts and consistent with solar abundances depending on the amount of oxygen locked up in CO. However, collisions can have caused a strong reduction of the water ice content.
总的来说,外卫星的水冰含量随位置和温度变化,不仅与距离太阳的距离有关,还与其母行星有关。在较冷的地区,冰含量通常较高(50%左右),并且与太阳丰度一致,取决于 CO 中氧的锁定量。然而,碰撞可能导致水冰含量大幅减少。

6.4 Water in the giant planets
巨大行星中的水

The largest reservoirs of what was once water ice in the solar nebula disk are presumably locked up in the giant planets. If the formation of giant planets started with an initial solid core of 10–15 Earth masses with subsequent growth from a swarm of planetesimals of ice and rock with solar composition, the giant planets could have had several Earth masses of oxygen, some or much of which may have been in water molecules in the original protoplanetary disk. The core also gravitationally attracts the surrounding gas in the disk consisting mostly of H and He with the other elements in solar composition. The resulting gas giant planet has a large mass and diameter, but a low overall density compared with rocky planets. The giant planet atmosphere is expected to have an excess in heavy elements, either due to the vaporization of the icy planetesimals when they entered the envelopes of the growing planet during the heating phase, or due to partial erosion of the original core, or both (Encrenaz, 2008; Mousis et al., 2009). These calculations assume that all heavy elements are equally trapped within the ices initially, which is a debatable assumption, and that the ices fully evaporate, with most of the refractory material sedimenting onto the core. In principle, the excesses provide insight into giant planet formation mechanisms and constraints on the composition of their building blocks.
太阳星云盘中曾经的水冰最大储量可能被锁定在巨大行星中。如果巨大行星的形成始于一个初始固态核心,质量为 10-15 个地球质量,随后从冰和岩石的行星体群中以太阳组成的方式增长,那么巨大行星可能拥有数个地球质量的氧气,其中一部分或大部分可能以水分子的形式存在于原始原行星盘中。核心还通过引力吸引盘中的周围气体,该气体主要由 H 和 He 组成,其他元素以太阳组成的方式存在。由此产生的气态巨行星拥有较大的质量和直径,但与岩石行星相比,整体密度较低。预计巨行星大气中会有过量的重元素,这可能是由于当它们在加热阶段进入不断增长的行星的包层时,冰质行星体蒸发,或者由于原始核心的部分侵蚀,或两者兼而有之。 这些计算假设所有重元素最初均被均匀地困在冰中,这是一个有争议的假设,并且这些冰完全蒸发,其中大部分不挥发物质沉积到了核心上。原则上,这些过剩物提供了对巨行星形成机制以及对它们构成基础的约束的见解。

6.4.1 Jupiter and Saturn  6.4.1 木星和土星

For Jupiter, elemental abundances can be derived from spectroscopic observations and from data collected by the Galileo probe, which descended into the Jovian atmosphere. Elements like C, N, S and the noble gases Ar, Kr and Xe, have measured excesses as expected at 4±\pm2 (Owen and Encrenaz, 2006). However, O appears to show significant depletion. Unfortunately, the deep oxygen abundance in Jupiter is not known. The measured abundance of gaseous water—the primary carrier of oxygen in the Jovian atmosphere since there is little CO—provides only a lower limit since the troposphere at about 100 mbar is a region of minimum temperature (\sim110 K for Jupiter) and therefore acts as a cold trap where water can freeze out. Thus, the amount of gaseous water is strongly affected by condensation and rainout associated with large-scale advective motions (Showman and Dowling, 2000), meteorological processes (Lunine and Hunten, 1987), or both. The Galileo probe fell into a so-called ‘hot spot’ (for the excess brightness observed in such regions at 5 μ\mum wavelengths), with enhanced transparency and hence depleted in water, and is thus not representative of the planet as a whole. The water abundance, less than 1/10 the solar value in the upper atmosphere, was observed to be higher at higher pressures, toward the end of the descent (Roos-Serote et al., 2004). The sparseness of the measurements made it impossible to know whether the water had ‘leveled out’ at a value corresponding to 1/3 solar or would have increased further had the probe returned data below the final 21 bar level.
对于木星,元素丰度可以从光谱观测和伽利略探测器收集的数据中推导出来,该探测器降入了木星大气层。像 C、N、S 和惰性气体 Ar、Kr 和 Xe 这样的元素,其过量测量值与预期相符(Owen 和 Encrenaz,2006)。然而,氧似乎显示出明显的耗尽。不幸的是,木星中的氧深度丰度尚不为人所知。测得的气态水丰度——作为木星大气中氧的主要载体,因为几乎没有 CO——仅提供一个下限,因为大气层在大约 100 毫巴处是一个最低温度区域(木星为 110K),因此起到了水可以结冰的冷陷阱作用。因此,气态水的量受到与大尺度平流运动相关的冷凝和降水以及气象过程(Showman 和 Dowling,2000),气象过程(Lunine 和 Hunten,1987)或两者的强烈影响。伽利略探测器坠入了所谓的“热点”(在 5 μ\mu m 波长处观察到这些区域的过量亮度),具有增强的透明度,因此水含量较低,并且因此不代表整个行星。 水的丰度,在上层大气中不到太阳值的十分之一,被观察到在较高压力下更高,在下降末端(Roos-Serote 等,2004 年)。测量的稀疏性使人无法知道水是否已经达到相当于太阳的三分之一的值,或者如果探测器返回了最终 21 巴以下的数据,水是否会进一步增加。

As noted above, predictions for standard models of planetesimal accretion—where volatiles are either adsorbed on, or enclathrated in, water ice—give oxygen abundances 3–10 times solar in the Jovian deep interior. Although atmospheric explanations for the depleted water abundance in Jupiter are attractive, one must not rule out the possibility that water truly is depleted in the Jovian interior—that is, the oxygen-to-hydrogen ratio in Jupiter is less enriched than the carbon value at 4±\pm2 times solar. A motivation for making such a case is that at least one planet with a C/O1—a ‘carbon-rich planet’—has been discovered >>(Madhusudhan et al., 2011a), companion to the star WASP12a with a C/O ratio of 0.44, roughly solar. One explanation is that the portion of this system’s protoplanetary disk was somehow depleted in water at the time the planet formed and acquired its heavy element inventory (Madhusudhan et al., 2011b).
如上所述,行星体积吸积的标准模型预测,挥发物要么被吸附在水冰上,要么被包裹在水冰中,会使木星深部的氧丰度是太阳的 3-10 倍。尽管大气解释对于木星中水的稀缺是有吸引力的,但不能排除水在木星内部真正稀缺的可能性,也就是说,木星中的氧氢比可能比碳值富集度低,只有太阳的 4 ±plus-or-minus\pm 2 倍。提出这种情况的动机之一是,至少发现了一个 C/O1 的行星——一个“碳丰富行星”(Madhusudhan 等人,2011a),它是 WASP12a 恒星的伴星,其 C/O 比为 0.44,大致等同于太阳。一个解释是,这个系统原行星盘的一部分在行星形成并获得重元素库存时,某种方式上水被耗尽了(Madhusudhan 等人,2011b)。

Prior to this discovery, the possibility of a carbon-rich Jupiter was considered on the basis of the Galileo results alone by Lodders (2004) who proposed that in the early solar system formation the snow line might have been further from the Sun than the point at which Jupiter formed, and volatiles adhering to solid organics rather than water ice were carried into Jupiter.
在这一发现之前,根据伽利略的结果,Lodders(2004)认为存在碳丰富的木星可能性,他提出在早期太阳系形成过程中,雪线可能比木星形成时距离太阳更远,并且挥发物质附着在固体有机物而非水冰上被带入木星。

Mousis et al. (2012) looked at the possibility that Jupiter may have acquired planetesimals from an oxygen-depleted region by examining element-by-element the fit to the Galileo probe data of two contrasting models: one in which the planetesimal building blocks of Jupiter derived from a disk with C/O= 1/2 (roughly, the solar value), and the other in which C/O=1. Within the curent error bars, the two cases cannot be distinguished. However, any determination of the deep oxygen abundance in Jupiter yielding an enrichment of less than 2 times solar would be a strong argument in favor of a water depletion, and thus high C/O ratio, at certain places and times in the solar nebula disk.
Mousis 等人(2012 年)研究了木星可能从氧贫区域获得小行星的可能性,通过逐个元素地检查两种对比模型与伽利略探测器数据的匹配情况:一种是木星的小行星构建模块源自 C/O=1/2(大致为太阳值)的盘,另一种是 C/O=1 的盘。在当前误差范围内,这两种情况无法区分。然而,任何确定木星深层氧丰度的结果,如果低于太阳值的 2 倍,将是支持水贫化和因此太阳星云盘中某些地点和时间高 C/O 比的有力论据。

Are such depletions plausible? A wide range of oxidation states existed in the solar nebula disk at different times and locations. For example, the driest rocks, the enstatite chondrites, are thought to have come from parent bodies formed inward of all the other parent bodies, and their mineralogy suggests reducing conditions—consistent with a depleted water vapor abundance—in the region of the nebula where they formed (Krot et al., 2000). One recent model of the early evolution of Jupiter and Saturn hypothesizes that these giant planets moved inward significantly during the late stages of their formation, reaching 1.5 AU in the case of Jupiter (Walsh et al., 2011), almost certainly inward of the snow line (see § 8.2). If planetesimals in this region accreted volatiles on refractory organic and silicate surfaces which were then incorporated into Jupiter, the latter would appear carbon-rich and oxygen-depleted. However, temperatures in that region may not have been low enough to provide sufficient amounts of the more volatile phases. A second possibility is that Jupiter’s migration scattered these water poor planetesimals to the colder outer solar system, where they trapped noble gases and carbon-and nitrogen-bearing species at lower temperature—but water ice, already frozen out, was not available.
这些减少是可信的吗?太阳星云盘在不同时间和地点存在着广泛的氧化状态。例如,最干燥的岩石——榍石球粒陨石,被认为来自形成在所有其他母体内部的母体,并且它们的矿物学表明在它们形成的星云区域存在还原条件——与缺乏水蒸气丰度一致(Krot 等人,2000 年)。最近有一个关于木星和土星早期演化的模型假设这些巨大行星在形成的后期阶段明显向内移动,例如,木星的情况下达到了 1.5 AU(Walsh 等人,2011 年),几乎可以确定在雪线内部(见§ 8.2)。如果该区域的小行星在难挥发性有机和硅酸盐表面上吸附挥发物质,然后这些物质被纳入木星,那么后者将呈现出富含碳且缺氧的特征。然而,该区域的温度可能不低到足以提供足够数量的更易挥发的相。 第二种可能性是,木星的迁移将这些贫水行星碎片散布到更冷的外太阳系,那里它们在较低温度下捕获了惰性气体和含碳和氮的物质,但已经冷冻结晶的水冰却不可用。

In order to test such models, one must directly determine the abundance of oxygen-bearing species in Jupiter and if possible, in Saturn. The case of Saturn is similar to that of Jupiter in the sense that the carbon excess as derived from CH4 spectroscopy is as expected, but again no reliable oxygen abundance can be determined. Under the conditions present in the Jovian envelope at least (if not its core), water will dominate regardless of the initial carbon oxidation state in the solar nebula. NASA’s Juno mission to Jupiter will measure the water abundance down to many tens of bars via a microwave radiometer (MWR) (Janssen et al., 2005). Complementary to the MWR is a near-infrared spectrometer JIRAM (Jovian Infrared Auroral Mapper), that will obtain the water abundance in the meteorological layer (Adriani et al., 2008). The two instruments together will be able to provide a definitive answer for whether the water abundance is below or above solar, and in the latter case, by how much, when Juno arrives in 2016. Juno will also determine the mass of the heavy element core of Jupiter, allowing for an interpretation of the significance of the envelope water abundance in terms of total oxygen inventory.
为了测试这些模型,人们必须直接确定木星和可能的土星中含氧物种的丰度。土星的情况与木星类似,从 CH 4 光谱推导出的碳过量与预期相符,但再次无法确定可靠的氧丰度。至少在木星包层中的条件下(如果不考虑其核心),无论太阳星云中的碳氧化状态如何,水都将占主导地位。美国宇航局的朱诺任务将通过微波辐射计(MWR)(Janssen 等人,2005 年)将水的丰度测量到许多巴。与 MWR 相辅相成的是近红外光谱仪 JIRAM(木星红外极光映像仪),它将获取气象层中的水的丰度(Adriani 等人,2008 年)。这两种仪器将能够在朱诺于 2016 年抵达时提供一个明确的答案,即水的丰度是低于还是高于太阳,以及在后一种情况下,超过多少。 朱诺还将确定木星的重元素核的质量,从而可以解释包裹水丰度在总氧库方面的意义。

There is no approved mission yet to send a probe into Saturn akin to Galileo. However, the Cassini Saturn Orbiter will make very close flybys of Saturn starting in 2016, similar to what Juno will do at Jupiter. Unfortunately, a microwave radiometer akin to that on Juno is not present on Cassini, but a determination of the heavy element core mass of Saturn may be obtained from remote sensing. In summary, the fascinating possibility that Jupiter and Saturn may have distinct oxygen abundances because they sampled at different times and to differing extents regions of the solar nebula disk heterogeneous in oxygen (i.e., water) abundance is testable if the bulk oxygen abundances can be measured.
目前尚无批准的任务类似于伽利略将探测器送入土星。然而,卡西尼土星轨道飞行器将从 2016 年开始非常接近地飞越土星,类似于朱诺在木星上所做的。不幸的是,卡西尼上没有类似于朱诺上的微波辐射计,但可以通过遥感获得土星重元素核质量的确定。总之,令人着迷的可能性是,由于它们在不同时间和不同程度上取样了太阳星云盘中氧(即水)丰度异质的区域,木星和土星可能具有不同的氧丰度,如果可以测量总氧丰度,则可以进行测试。

6.4.2 Uranus and Neptune  6.4.2 天王星和海王星

The carbon excesses measured from CH4 near-infrared spectroscopy give values of 30–50 for Uranus and Neptune, compared with 9 and 4 for Saturn and Jupiter (Encrenaz, 2008). These values are consistent with the assumption of an initial core of 10–15 Earth masses with heavy elements in solar abundances. Unfortunately, nothing is known about the water or oxygen content in Uranus and Neptune, since water condensation occurs at such deep levels that even tropospheric water vapor cannot be detected. Unexpectedly, ISO detected water vapor in the upper atmospheres of both ice giants, with mixing ratios orders of magnitude higher than the saturation level at the temperature inversion (Feuchtgruber et al., 1997). These observations can only be accounted for by an external flux of water molecules, due to interplanetary dust, or sputtering from rings or satellites. Similarly, Herschel maps of water of Jupiter demonstrate that even 15 years after the Shoemaker-Levy 9 impact, more than 95% of the stratospheric Jovian water comes from the impact (Cavalié et al., 2013). On the other hand, the low D/H ratios in molecular hydrogen measured by Herschel in the atmospheres of Uranus and Neptune may imply a lower ice mass fraction of their cores than previously thought (14–32%) (Feuchtgruber et al., 2013) .
通过近红外光谱测量的碳过量值显示,与土星和木星的 9 和 4 相比,天王星和海王星的值为 30-50(Encrenaz,2008)。这些数值与初始核心质量为 10-15 个地球质量,含有太阳丰度的重元素的假设一致。不幸的是,由于水凝结发生在如此深的层次,甚至对流层水蒸气也无法被检测,因此对天王星和海王星的水或氧含量一无所知。出乎意料的是,ISO 在两颗冰巨星的上层大气中检测到了水蒸气,混合比高出温度反转点的饱和水平几个数量级(Feuchtgruber 等,1997)。这些观测只能通过外部水分子通量来解释,这是由于行星间尘埃,或者来自环或卫星的溅射。同样,对木星水的赫歇尔地图表明,即使在舒梅克-莱维 9 号撞击 15 年后,95%以上的平流层木星水来自撞击(Cavalié等,2013)。 另一方面,由 Herschel 在天王星和海王星大气中测量的分子氢中低的 D/H 比可能意味着它们的核心冰质量分数比先前认为的要低(14-32%)(Feuchtgruber 等,2013 年)。

It is unlikely that information on the deep oxygen abundance will be available for Uranus and Neptune anytime soon, since reaching below the upper layers to determine the bulk water abundance is very difficult. Thus, data complementary to that for Jupiter and Saturn will likely come first from observations of Neptune-like exoplanets.
由于很难达到深层以确定天王星和海王星的总水含量,因此在不久的将来可能不太可能获得有关深层氧丰度的信息。因此,与木星和土星的数据相辅相成的数据可能首先来自类海王星的系外行星的观测。

Refer to caption
Figure 10: Water content of various parent bodies at the time of Earth’s growth as a function of radial distance from the Sun. Enstatite chondrites originating from asteroids around 1.8 AU in the inner disk are very dry. In contrast, carbonaceous chondrites originating from the outer asteroid belt and beyond have a water content of 5–10%, and main belt comets, TNOs and regular comets even more. Figure by M. Persson, adapted from Morbidelli et al. (2012).
图 10:地球形成时各母天体的水含量与距离太阳的径向距离的关系。源自内部盘围绕 1.8 AU 的小行星的蛇纹石球粒陨石非常干燥。相比之下,源自外小行星带及更远处的碳质球粒陨石的水含量为 5-10%,而主带彗星、TNOs 和常规彗星甚至更高。图由 M. Persson 绘制,改编自 Morbidelli 等(2012 年)。

7 WATER IN THE INNER SOLAR SYSTEM
7 内部太阳系中的水

Terrestrial planets were formed by accretion of rocky planetesimals, and they have atmospheres that are only a small fraction of their total masses. They are small and have high densities compared with giant planets. Earth and Venus have very similar sizes and masses, whereas Mars has a mass of only 10% of that of Earth and a somewhat lower density (3.9 vs 5.2 gr cm-3). Terrestrial planets undergo different evolutionary processes compared with their gas-rich counterparts. In particular, their atmospheres result mostly from outgassing and from external bombardment. Internal differentiation of the solid material after formation leads to a structure in which the heavier elements sink to the center, resulting in a molten core (consisting of metals like iron and nickel), a mantle (consisting of a viscous hot dense layer of magnesium rich silicates) and a thin upper crust (consisting of colder rocks).
地球型行星是由岩石小行星的聚集形成的,它们的大气层仅占其总质量的一小部分。与巨大行星相比,它们体积小,密度高。地球和金星的大小和质量非常相似,而火星的质量仅为地球的 10%,密度略低(3.9 与 5.2 克/立方厘米)。地球型行星经历与富气体对应物不同的演化过程。特别是,它们的大气主要是由排气和外部轰击造成的。形成后,固体材料的内部分化导致了一种结构,其中较重的元素沉积到中心,形成了一个熔融的核心(由铁和镍等金属组成)、一个地幔(由富镁硅酸盐的粘稠热密层组成)和一个薄的上地壳(由较冷的岩石组成)。

Water has very different appearances on Venus, Earth and Mars, which is directly related to their distances from the Sun at 0.7, 1.0 and 1.7 AU, respectively. On Venus, with surface temperatures around 730 K, only gaseous water is found, whereas the surface of Mars has seasonal variations ranging from 150–300 K resulting in water freezing and sublimation. Its current mean surface pressure of 6 mbar is too low for liquid water to exist, but there is ample evidence for liquid water on Mars in its early history. Most of the water currently on Mars is likely subsurface in the crust down to 2 km depth. Earth is unique in that its mean surface temperature of 288 K and pressure of 1 bar allow all three forms of water to be present: vapor, liquid and ice. In § 8, the origin of water on these three planets will be further discussed. It is important to keep in mind that even though these planets have very different atmospheres today, they may well have started out with comparable initial water mass fractions and similar atmospheric compositions dominated by H2O, CO2, and N2 (Encrenaz, 2008).
金星、地球和火星上的水呈现出非常不同的外观,这与它们分别距离太阳的距离有直接关系,分别为 0.7、1.0 和 1.7 天文单位。在金星上,表面温度约为 730K,只发现了水蒸气,而火星表面的季节变化范围在 150-300K 之间,导致水结冰和升华。火星目前的平均表面压力为 6 毫巴,太低以至于无法存在液态水,但有充分证据表明火星早期存在液态水。目前火星上大部分水很可能位于地壳下达 2 公里深处。地球独特之处在于其平均表面温度为 288K,压力为 1 大气压,使得水的三种形态都可以存在:水蒸气、液态和冰。在第 8 节中,将进一步讨论这三个行星上水的起源。重要的是要记住,即使这些行星今天的大气非常不同,它们很可能最初具有相当的初始水质量分数和由 H 2 O、CO 2 和 N 2 主导的类似大气组成(Encrenaz,2008)。

Earth has a current water content that is non-negligible. The mass of the water contained in the Earth’s crust (including the oceans and the atmosphere) is 2.8×1042.8\times 10^{-4} Earth masses, denoted as ‘one Earth ocean’ because almost all of this is in the surface waters of the Earth. The mass of the water in the present-day mantle is uncertain. Lécuyer et al. (1998) estimate it to be in the range of (0.8–8)×104\times 10^{-4} Earth masses, equivalent to 0.3–3 Earth oceans. More recently, Marty (2012) provides arguments in favor of a mantle water content as high as \sim7 Earth oceans. However, an even larger quantity of water may have resided in the primitive Earth and been subsequently lost during differentiation and impacts. Thus, the current Earth has a water content of roughly 0.1% by mass, larger than that of enstatite chondrites, and it is possible that the primitive Earth had a water content comparable with or larger than that of ordinary chondrites, definitely larger than the water content of meteorites and material that condensed at 1 AU (Fig. 10).
地球目前的水含量是不可忽略的。地球地壳中含有的水的质量(包括海洋和大气层)为 2.8×1042.8superscript1042.8\times 10^{-4} 地球质量,被称为“一个地球海洋”,因为几乎所有这些水都在地球的表层水中。目前地幔中的水的质量是不确定的。Lécuyer 等人(1998 年)估计其在(0.8-8) ×104absentsuperscript104\times 10^{-4} 地球质量的范围内,相当于 0.3-3 个地球海洋。最近,Marty(2012 年)提出了支持地幔水含量高达 similar-to\sim 7 个地球海洋的论点。然而,更多的水可能曾经存在于原始地球中,并在分化和撞击过程中被丢失。因此,目前的地球的水含量大约占质量的 0.1%,比榴石隕石的水含量更高,可能原始地球的水含量与普通隕石相当或更高,绝对比隕石和在 1 AU 处凝结的物质的水含量更高(图 10)。

Remote sensing and space probes have detected water in the lower atmosphere of Venus at a level of a few tens of parts per millions (Gurwell et al., 2007), confirming that it is a minor component. Radar images show that the surface of Venus is covered with volcanoes and is likely very young, only 1 billion years. Direct sampling of the Venusian crust is needed to determine the extent of past and present hydration which will provide an indication of whether Venus once had an amount of water comparable to that on Earth. Hints that there may have been much more water on Venus and Mars come from D/H measurements (see § 8.1).
遥感和空间探测器已经在金星的下层大气中探测到水,水平为每百万分之几十(Gurwell 等人,2007),证实它是一个次要成分。雷达图像显示,金星表面覆盖着火山,并且可能非常年轻,只有 10 亿年。需要直接取样金星地壳以确定过去和现在的水合程度,这将表明金星曾经是否有与地球上相当的水量。来自 D/H 测量的线索表明金星和火星上可能曾经有更多的水(见§ 8.1)。

8 ORIGIN OF WATER TERRESTRIAL PLANETS
8 水的起源 地球类行星

Earth, Mars and perhaps also Venus all have water mass fractions that are higher than those found in meteorites in the inner solar nebula disk. Where did this water come from, if the local planetesimals were dry? There are two lines of arguments that provide clues to its origin. One clue comes from measuring the D/H ratio of the various water mass reservoirs. The second argument looks at the water mass fractions of various types of asteroids and comets, combined with the mass delivery rates of these objects on the young planets at the time of their formation, based on models of the dynamics of these planetesimals. For example, comets have plenty of water, but the number of comets that enter the inner 2 AU and collide with proto-Earth or -Mars is small. To get one Earth ocean, 108\sim 10^{8} impacting comets are needed. It is therefore not a priori obvious which of the reservoirs shown in Fig. 10 dominates.
地球、火星,也许还有金星,它们的水质量分数都高于内部太阳星云盘中陨石中发现的水质量分数。如果当地的小行星干燥,那么这些水是从哪里来的呢?有两条线索可以提供关于其起源的线索。一条线索来自于测量各种水质量储集库的 D/H 比。第二个论点则关注各种类型小行星和彗星的水质量分数,结合这些物体在年轻行星形成时的质量传递速率,基于这些小行星动力学模型。例如,彗星有大量水,但进入内部 2 AU 并与原始地球或火星碰撞的彗星数量很少。要得到一个地球海洋,需要撞击的彗星数量为 0。因此,显然不明显图 10 中显示的哪个储集库占主导地位。

Related to this point is the question whether the planets formed ‘wet’ or ‘dry’. In the wet scenario, the planets either accreted a water-rich atmosphere or they formed from planetesimals with water bonded to silicate grains at 1 AU, with subsequent outgassing of water as the material is heated up during planet formation. In the dry scenario, the terrestrial planets are initially built up from planetesimals with low water mass fractions and water is delivered to their surfaces by water-rich planetesimals, either late in the building phase, or even later after the planets have formed and differentiated. If after differentiation, this scenario is often called a ‘late veneer’. In the following, we first discuss the D/H ratios and then come back to the mass fraction arguments.
与这一点相关的问题是行星是“湿润”还是“干燥”形成的。在湿润的情况下,行星要么吸积了富含水的大气层,要么是从 1 AU 处带有水与硅酸盐颗粒结合的小行星形成的,随后在行星形成过程中加热时释放出水。在干燥的情况下,类地行星最初是由低水质量分数的小行星构建起来的,并且水是通过富含水的小行星向它们的表面输送的,要么是在建造阶段的后期,要么是在行星形成并分化后的更晚阶段。如果在分化后,这种情况通常被称为“晚期外壳”。在接下来的内容中,我们首先讨论 D/H 比率,然后再回到质量分数的论点。

8.1 D/H ratios in solar system water reservoirs
太阳系水库中的 D/H 比率

The D/H ratio of water potentially provides a unique fingerprint of the origin and thermal history of water. Figure 11 summarizes the various measurements of solar system bodies, as well as interstellar ices and protostellar objects.
水的 D/H 比率可能提供了水的起源和热史的独特指纹。图 11 总结了太阳系天体以及星际冰和原恒星物体的各种测量结果。

Refer to caption
Figure 11: D/H ratio in water in comets and warm protostellar envelopes compared to values in the Earth oceans, the giant planets, the solar nebula disk, and the interstellar medium. Values for carbonaceous meteorites (CI), the Moon and Saturn’s moon Enceladus are presented as well. Note that both (D/H)X, the deuterium to hydrogen ratio in molecule X, and HDO/H2O are plotted; measured HD/H2 and HDO/H2O ratios are 2×\times(D/H)X. Figure by M. Persson, based on Bockelée-Morvan et al. (2012) and subsequent measurements cited in the text. The protostellar data refer to warm gas (Persson et al., 2013, and subm.)
图 11:彗星和温暖原恒星包层中水的 D/H 比值与地球海洋、巨大行星、太阳星云盘和星际介质中的数值进行比较。碳质陨石(CI)、月球和土星的卫星土卫二的数值也呈现出来。请注意,图中绘制了分子 X 中的(D/H) X ,以及 HDO/H 2 O;测得的 HD/H 2 和 HDO/H 2 O 比值为 2 ×\times (D/H) X 。图由 M. Persson 制作,基于 Bockelée-Morvan 等人(2012)和文中引用的后续测量数据。原恒星数据指的是温暖气体(Persson 等人,2013 年,及后续)。

The primordial [D]/[H] ratio set shortly after the Big Bang is 2.7×1052.7\times 10^{-5}. Since then, deuterium has been lost due to nuclear fusion in stars, so the deuterium abundance at the time of our solar system formation, 4.6 billion years ago (redshift of about z1.4z\approx 1.4), should be lower than the primordial value, but higher than the current day interstellar [D]/[H] ratio. The latter value has been measured in the diffuse local interstellar medium from UV absorption lines of atomic D and H, and varies from place to place but can be as high as 2.3×1052.3\times 10^{-5} (Prodanović et al., 2010). These higher values suggest that relatively little deuterium has been converted in stars, so a [D]/[H] value of the solar nebula disk only slightly above the current ISM value is expected at the time of solar system formation. The solar nebula [D]/[H] ratio can be measured from the solar wind composition as well as from HD/H2 in the atmospheres of Jupiter and Saturn and is found to be (2.5±0.5)×105(2.5\pm 0.5)\times 10^{-5} (Robert et al., 2000, and references cited), indeed in-between the primordial and the current ISM ratios.444Note that measured HD/H2 and HDO/H2O ratios are 2×\times(D/H)X, the D/H ratios in these molecules.
在大爆炸后不久设定的原始[D]/[H]比例为 2.7×1052.7superscript1052.7\times 10^{-5} 。此后,氘在恒星核聚变中丢失,因此在我们太阳系形成时,即 46 亿年前(红移约 z1.41.4z\approx 1.4 ),氘的丰度应低于原始值,但高于当前星际[D]/[H]比值。后者的数值已通过原子 D 和 H 的紫外吸收线在弥散的局部星际介质中测量,不同地方有所变化,但可以高达 2.3×1052.3superscript1052.3\times 10^{-5} (Prodanović等,2010)。这些较高的数值表明相对较少的氘已在恒星中转化,因此预计在太阳系形成时,太阳星云盘的[D]/[H]值仅略高于当前 ISM 值。太阳星云[D]/[H]比值可以从太阳风组成以及木星和土星大气中的 HD/H 2 中测量,发现为 (2.5±0.5)×105plus-or-minus2.50.5superscript105(2.5\pm 0.5)\times 10^{-5} (Robert 等,2000 年,以及引用的参考文献),确实介于原始值和当前 ISM 比值之间。 444Note that measured HD/H2 and HDO/H2O ratios are 2×\times(D/H)X, the D/H ratios in these molecules.

The D/H ratio of Earth’s ocean water is 1.5576×1041.5576\times 10^{-4} (‘Vienna Standard Mean Ocean Water’, VSMOW or SMOW). Whether this value is representative of the bulk of Earth’s water remains unclear, as no measurements exist for the mantle or the core. It is thought that recycling of water in the deep mantle does not significantly change the D/H ratio. In any case, the water D/H ratio is at least a factor of 6 higher than that of the gas out of which our solar system formed. Thus, Earth’s water must have undergone fractionation processes that enhance deuterium relative to hydrogen at some stages during its history, of the kind described in § 2.4 for interstellar chemistry at low temperatures.
地球海洋水的 D/H 比例为 1.5576×1041.5576superscript1041.5576\times 10^{-4} ('维也纳标准海洋水',VSMOW 或 SMOW)。地球水的大部分是否代表这个值仍不清楚,因为地幔或地核没有测量数据。人们认为,深部地幔中的水循环不会显著改变 D/H 比例。无论如何,水的 D/H 比例至少比太阳系形成时的气体高出 6 倍。因此,地球的水在其历史某些阶段必须经历了增强重氢相对于氢的分馏过程,这类过程在低温下的星际化学中描述在第 2.4 节。

Which solar system bodies show D/H ratios in water similar to those found in Earth’s oceans? The highest [D]/[H] ratios are found in two types of primitive meteorites (Robert et al., 2000): LL3, an ordinary chondrite that may have the near-Earth asteroid 433 Eros as its parent body (up to a factor of 44 enhancement compared to the protosolar ratio in H2) and some carbonaceous chondrites (a factor of 15 to 25 enhancement). Most chondrites show lower enhancements than the most primitive meteorites. Pre-Herschel observations of six Oort-cloud or long-period comets give a D/H ratio in water of 3×104\sim 3\times 10^{-4}, a factor of 12 higher than the protosolar ratio in H2 (Mumma and Charnley, 2011).
哪些太阳系天体显示出与地球海洋中发现的水的 D/H 比类似的比例?最高的[D]/[H]比例出现在两种原始陨石中(Robert 等人,2000 年):LL3,一种可能具有近地小行星 433 Eros 作为其母体的普通球粒陨石(与 H 中的原始太阳比例相比增加了 44 倍),以及一些含碳球粒陨石(增加了 15 到 25 倍)。大多数球粒陨石显示出比最原始的陨石更低的增强效应。对六颗奥尔特云或长周期彗星的赫歇尔之前的观测显示,水的 D/H 比例为 1,比 H 中的原始太阳比例高 12 倍(Mumma 和 Charnley,2011)。

Herschel provided the first measurement of the D/H ratio in a Jupiter-family comet. A low value of (1.61±0.24)×104(1.61\pm 0.24)\times 10^{-4}, consistent with VSMOW, was measured in comet 103P/Hartley 2 (Hartogh et al., 2011). In additon, a relatively low ratio of (2.06±0.22)×104(2.06\pm 0.22)\times 10^{-4} was found in the Oort-cloud comet C/2009 P1 Garradd (Bockelée-Morvan et al., 2012) and a sensitive upper limit of <2×104<2\times 10^{-4} (3σ\sigma) was obtained in another Jupiter-family comet 45P/Honda-Mrkos-Pajdušáková (Lis et al. 2013). The Jupiter-family comets 103P and 45P are thought to originate from the large reservoir of water-rich material in the Kuiper belt or scattered disk at 30–50 AU. In contrast, Oort cloud comets, currently at much larger distances from the Sun, may have formed closer in, near the current orbits of the giant planets at 5–20 AU, although this traditional view has been challened in recent years (see § 6.2). Another caveat is that isotopic ratios in cometary water may have been altered by the outgassing process (Brown et al., 2011). Either way, the Herschel observations demonstrate that the earlier high D/H values are not representative of all comets.
赫歇尔首次测量了木星家族彗星中的 D/H 比值。在哈特利 2 号彗星(Hartogh 等人,2011)中测得了与 VSMOW 一致的低值 (1.61±0.24)×104plus-or-minus1.610.24superscript104(1.61\pm 0.24)\times 10^{-4} 。此外,在奥尔特云彗星 C/2009 P1 Garradd(Bockelée-Morvan 等人,2012)中发现了相对较低的比值 (2.06±0.22)×104plus-or-minus2.060.22superscript104(2.06\pm 0.22)\times 10^{-4} ,并在另一颗木星家族彗星 45P/Honda-Mrkos-Pajdušáková(Lis 等人,2013)中获得了 <2×104absent2superscript104<2\times 10^{-4} (3 σ\sigma )的敏感上限。据认为,103P 和 45P 木星家族彗星起源于库伯带或分散盘中富含水的大量物质,距太阳 30-50 天文单位。相比之下,奥尔特云彗星目前距离太阳更远,可能在更接近的地方形成,靠近 5-20 天文单位的巨行星当前轨道,尽管这种传统观点近年来受到挑战(见第 6.2 节)。另一个警告是,彗星水中的同位素比值可能已经被气体抛射过程改变(Brown 等人,2011)。无论如何,赫歇尔的观测表明,早期高 D/H 值并不代表所有彗星。

How do these solar system values compare with interstellar and protostellar D/H ratios? Measured D/H ratios in water in protostellar envelopes vary strongly. In the cold outer envelope, D/H values for gas-phase water are as high as 10210^{-2} (Liu et al., 2011; Coutens et al., 2012, 2013), larger than the upper limits in ices of (2–5)<<×103\times 10^{-3} obtained from infrared spectroscopy (Dartois et al., 2003; Parise et al., 2003). In the inner warm envelope, previous discrepancies for gaseous water appear to have been resolved in favor of the lower values, down to (3–5)×104\times 10^{-4} (Jørgensen and van Dishoeck, 2010a; Visser et al., 2013; Persson et al., 2013, and subm.) The latter values are within a factor of two of the cometary values.
这些太阳系数值与星际和原恒星 D/H 比值相比如何?在原恒星包层中,水的 D/H 比值变化很大。在寒冷的外包层中,气相水的 D/H 值高达 102superscript10210^{-2} (Liu 等,2011 年;Coutens 等,2012 年,2013 年),比从红外光谱中获得的冰的上限值(2-5) << ×103absentsuperscript103\times 10^{-3} 要大(Dartois 等,2003 年;Parise 等,2003 年)。在内部温暖的包层中,以前对气态水的差异似乎已经解决,有利于更低的值,降至(3-5) ×104absentsuperscript104\times 10^{-4} (Jørgensen 和 van Dishoeck,2010a;Visser 等,2013 年;Persson 等,2013 年,及提交)。后者的值在彗星值的两倍范围内。

Based on these data, the jury is still out whether the D/H ratio in solar system water was already set by ices in the early (pre-)collapse phase and transported largely unaltered to the comet-forming zone (cf. Fig. 4), or whether further alteration of D/H took place in the solar nebula disk along the lines described in § 2.4.5 (see also chapter by Ceccarelli et al.). The original D/H ratio in ices may even have been reset early in the embedded phase by thermal cycling of material due to accretion events onto the star (see chapter by Audard et al.). Is the high value of 10210^{-2} found in cold gas preserved in the material entering the disk or are the lower values of <103<10^{-3} found in hot cores and ices more representative? Or do the different values reflect different D/H ratios in layered ices (Taquet et al., 2013)?
根据这些数据,关于太阳系水中的 D/H 比是否已经由早期(前)坍缩阶段的冰设置,并且大部分未经改变地运输到形成彗星的区域(参见图 4),目前尚无定论,或者 D/H 的进一步改变是否发生在太阳星云盘中,沿着§2.4.5 描述的方式进行(另请参阅 Ceccarelli 等人的章节)。冰中的原始 D/H 比甚至可能在嵌入阶段早期被由于物质对恒星的吸积事件而导致的热循环重置(请参阅 Audard 等人的章节)。在进入盘中的材料中保留的冷气体中发现的 102superscript10210^{-2} 的高值是存在的,还是在热核和冰中发现的 <103absentsuperscript103<10^{-3} 的较低值更具代表性?或者不同的值是否反映了分层冰中不同的 D/H 比(Taquet 等人,2013 年)?

Regardless of the precise initial value, models have shown that vertical and radial mixing within the solar nebula disk reduces the D/H ratios from initial values as high as 10210^{-2} to values as low as 10410^{-4} in the comet-forming zones (Willacy and Woods, 2009; Yang et al., 2012b; Jacquet and Robert, 2013; Furuya et al., 2013; Albertsson et al., 2014). For water on Earth, the fact that the D enrichment is a factor \sim6 above the solar nebula value rules out the warm thermal exchange reaction of H2O + HD \to HDO + H2 in the inner nebula (\sim1 AU) as the sole cause. Mixing with some cold reservoir with enhanced D/H at larger distances is needed.
无论初始值是多少,模型表明,太阳星云盘内的垂直和径向混合将将 D/H 比值从最初高达 102superscript10210^{-2} 的值降低到彗星形成区域低至 104superscript10410^{-4} 的值(Willacy 和 Woods,2009 年;Yang 等,2012b 年;Jacquet 和 Robert,2013 年;Furuya 等,2013 年;Albertsson 等,2014 年)。对于地球上的水来说,D 富集是太阳星云值的 similar-to\sim 6 倍,排除了 H 2 O + HD \to HDO + H 2 在内部星云( similar-to\sim 1 AU)中作为唯一原因的温暖热交换反应。需要与一些在较远距离具有增强 D/H 的冷贮库混合。

Figure 11 contains values for several other solar system targets. Enceladus, one of Saturn’s moons with volcanic activity, has a high D/H ratio consistent with it being built up from outer solar system planetesimals. Some measurements indicate that lunar water may have a factor of two higher hydrogen isotopic ratio than the Earth’s oceans (Greenwood et al., 2011), although this has been refuted (Saal et al., 2013).
图 11 包含了几个其他太阳系目标的数值。土卫二的土卫星之一,具有火山活动的土卫二,具有与由外太阳系小行星堆积而成的高 D/H 比值一致。一些测量表明,月球水的氢同位素比地球海洋高出两倍,尽管这一观点已被驳斥(Greenwood 等,2011 年;Saal 等,2013 年)。

Mars and Venus have interestingly high D/H ratios. The D/H ratio of water measured in the Martian atmosphere is 5.5 times VSMOW, which is interpreted to imply a significant loss of water over Martian history and associated enhancement of deuterated water. Because D is heavier than H, it escapes more slowly, therefore over time the atmosphere is enriched in deuterated species like HDO. A time-dependent model for the enrichment of deuterium on Mars assuming a rough outgassing efficiency of 50% suggests that the amount of water that Mars must have accreted is 0.04–0.4 oceans. The amount of outgassing may be tested by Curiosity Mars rover measurements of other isotopic ratios, such as 38Ar to 36Ar.
火星和金星的 D/H 比率非常高。在火星大气中测得的水的 D/H 比率是 VSMOW 的 5.5 倍,这被解释为意味着火星历史上水的显著流失和氘水的增加。因为 D 比 H 重,所以它逃逸得更慢,因此随着时间的推移,大气中富集了 HDO 等氘化物种。假设火星上氘的富集是一个时间相关模型,假设粗略的气体抽出效率为 50%,这表明火星必须吸收的水量为 0.04-0.4 个海洋。气体抽出的量可以通过好奇号火星探测器测量其他同位素比率来测试,例如 38 Ar 到 36 Ar。

The extremely high D/H ratio for water measured in the atmosphere of Venus, about 100 times VSMOW, is almost certainly a result of early loss of substantial amounts of crustal waters, regardless of whether the starting value was equal to VSMOW or twice that value. The amount of water lost is very uncertain because the mechanism and rate of loss affects the deuterium fractionation (Donahue and Hodges, 1992). Solar wind stripping, hydrodynamic escape, and thermal (Jeans’) escape have very different efficiencies for the enrichment of deuterium per unit amount of water lost. Values ranging from as little as 0.1% to one Earth ocean have been proposed.
金星大气中水的 D/H 比值极高,约为 VSMOW 的 100 倍,几乎可以肯定是由于大量地壳水的早期流失所致,无论起始值是否等于 VSMOW 或是其两倍。流失的水量非常不确定,因为流失的机制和速率会影响氘分馏(Donahue 和 Hodges,1992)。太阳风剥离、流体动力逃逸和热(James')逃逸对单位流失水量的氘富集效率差异很大。已经提出了从仅有 0.1%到一个地球海洋的数值范围。

8.2 Water delivery to the terrestrial planet zone
8.2 地球类行星区的水输送

8.2.1 Dry scenario  8.2.1 干旱情景

The above summary indicates at least two reservoirs of water ice rich material in the present-day solar system with D/H ratios consistent with that in Earth’s oceans: the outer asteroid belt and the Kuiper belt or scattered disk (as traced by comets 103P and 45P). As illustrated in Fig. 10, these same reservoirs also have a high enough water mass fraction to deliver the overall water content of terrestrial planets. The key question is then whether the delivery of water from these reservoirs is consistent with the current understanding of the early solar system dynamics.
上述摘要表明,至少在当前太阳系中存在至少两个富含水冰材料的水库,其 D/H 比与地球海洋中的比例一致:外部小行星带和柯伊伯带或散布盘(由彗星 103P 和 45P 追踪)。正如图 10 所示,这些相同的水库也具有足够高的水质量分数,可以输送地球类行星的整体含水量。关键问题是,来自这些水库的水的输送是否与当前对早期太阳系动力学的理解一致。

Over the last decade, there has been increasing evidence that the giant planets did not form and stay at their current location in the solar system but migrated through the protosolar disk. The Grand Tack scenario (Walsh et al., 2011) invokes movement of Jupiter just after it formed in the early, gas-rich stage of the disk (few Myr). The Nice model describes the dynamics and migration of the giant planets in the much later, gas-poor phase of the disk, some 800–900 Myr after formation (Gomes et al., 2005; Morbidelli et al., 2005; Tsiganis et al., 2005), see chapter by Raymond et al.
在过去的十年里,越来越多的证据表明,巨大行星并没有在太阳系中形成并停留在它们当前的位置,而是通过原太阳盘迁移。Grand Tack 情景(Walsh 等,2011)援引了木星在它形成后的早期、富气体阶段的盘中移动(几百万年)。尼斯模型描述了巨大行星在盘的后期、缺乏气体的阶段中的动力学和迁移,形成后大约 800-900 百万年(Gomes 等,2005;Morbidelli 等,2005;Tsiganis 等,2005),请参见 Raymond 等人的章节。

Refer to caption
Figure 12: Cartoon of the delivery of water-rich planetesimals in the outer asteroid belt and terrestrial planet-forming zone based on the Grand Tack scenario in which Jupiter and Saturn first move inwards to 1.5 AU and then back out again. Figure based on Walsh et al. (2011).
图 12:根据大回旋情景,展示了水富含的小行星在外部小行星带和类地行星形成区域的运送过程,其中木星和土星首先向内移动到 1.5 天文单位,然后再次向外移动。图基于 Walsh 等人(2011)的研究。

The Grand Tack model posits that if Jupiter formed earlier than Saturn just outside the snow line, both planets would have first migrated inward until they got locked in the 3:2 resonance, then they would have migrated outward together until the complete disappearance of the gas disk (see cartoon in Fig. 12). As summarized by Morbidelli et al. (2012), the reversal of Jupiter’s migration at 1.5 AU provides a natural explanation for the outer edge of the inner disk of (dry) embryos and planetesimals at 1 AU, which is required to explain the low mass of Mars. The asteroid belt at 2–4 AU is first fully depleted and then repopulated by a small fraction of the planetesimals scattered by the giant planets during their formation. The outer part of this belt around 4 AU would have been mainly populated by planetesimals originally beyond the orbits of the giant planets (5 AU), which explains similarities between primitive asteroids (C-type) and comets. Simulations by >>Walsh et al. (2011) show that as the outer asteroid belt is repopulated, (311)×102(3-11)\times 10^{-2} Earth masses of C-type material enters the terrestrial planet region, which exceeds by a factor 6–22 the minimum mass required to bring the current amount of water to the Earth. In this picture, Kuiper belt objects would at best be only minor contributors to the Earth’s water budget.
大规模调查模型认为,如果木星比土星早在雪线外形成,那么这两颗行星首先会向内迁移,直到它们被锁定在 3:2 的共振中,然后它们会一起向外迁移,直到气体盘完全消失。根据 Morbidelli 等人(2012)的总结,木星在 1.5 天文单位处的迁移逆转为(干燥)胚胎和行星碎片盘的内部边缘提供了一个自然解释,这是解释火星低质量所必需的。位于 2-4 天文单位处的小行星带首先完全耗尽,然后由巨大行星在形成过程中散射的一小部分行星碎片重新填充。这个小行星带的外部部分大约在 4 天文单位处主要由最初位于巨大行星轨道之外(5 天文单位)的行星碎片填充,这解释了原始小行星(C 型)和彗星之间的相似性。Walsh 等人的模拟。 (2011) show that as the outer asteroid belt is repopulated, (311)×102311superscript102(3-11)\times 10^{-2} Earth masses of C-type material enters the terrestrial planet region, which exceeds by a factor 6–22 the minimum mass required to bring the current amount of water to the Earth. In this picture, Kuiper belt objects would at best be only minor contributors to the Earth’s water budget.

The Nice model identifies Jupiter and Saturn crossing their 1:2 orbital resonance as the next key event in the dynamical evolution of the disk in the gas-poor phase. After the resonance crossing time at 880\sim 880 Myr, the orbits of the ice giants, Uranus and Neptune, became unstable. They then disrupted the outer disk and scattered objects throughout the solar system, including into the terrestrial planet region. However, Gomes et al. (2005) estimate the amount of cometary material delivered to the Earth to be only about 6% of the current ocean mass. A larger influx of material from the asteroid belt is expected, as resonances between the orbits of asteroids and giant planets can drive objects onto orbits with eccentricities and inclinations large enough to allow them to evolve into the inner solar system. By this time, the terrestrial planets should have formed already, so this material would be part of a ‘late delivery’.
尼斯模型确定了木星和土星跨越它们的 1:2 轨道共振作为气体贫乏阶段盘的动力学演化中的下一个关键事件。在 880similar-toabsent880\sim 880 百万年的共振穿越时间之后,冰巨星天王星和海王星的轨道变得不稳定。它们随后扰乱了外部盘,并将物体散布到整个太阳系,包括进入类地行星区域。然而,Gomes 等人(2005 年)估计传送到地球的彗星物质量仅约为当前海洋质量的 6%。预计将有更多来自小行星带的物质涌入,因为小行星和巨行星轨道之间的共振可以将物体驱动到具有足够大离心率和倾角的轨道,使它们能够演化到内太阳系。到那时,类地行星应该已经形成,因此这些物质将成为“后期交付”的一部分。

The Grand Tack model provides an attractive, although not unchallenged, explanation for the observed morphology of the inner solar system and the delivery of water to the Earth. In the asteroidal scenario, water is accreted during the formation phase of the terrestrial planets and not afterwards through bombardment as a late veneer. Typically, 50% of the water is accreted after the Earth has reached 60–70% of its final mass (Morbidelli et al., 2012). This appears in contradiction with the measurements of the distinct D/H ratio in lunar water mentioned above, approximately twice that in the Earth’s oceans (Greenwood et al., 2011). If valid, this would indicate that a significant delivery of cometary water to the Earth-Moon system occurred shortly after the Moon-forming impact. The Earth water would thus be a late addition, resulting from only one, or at most a few collisions with the Earth that missed the Moon (Robert, 2011). Expanding the sample of objects with accurate D/H measurements is thus a high priority, long-term science goal for the new submillimeter facilities.
Grand Tack 模型提供了一个有吸引力的,尽管并非毫无挑战的解释,用于解释内部太阳系的观察形态以及将水输送到地球。在小行星场景中,水在地球类行星形成阶段被吸积,而不是后来通过轰击作为晚期表面层。通常,在地球达到其最终质量的 60-70%后,50%的水被吸积(Morbidelli 等人,2012)。这似乎与上述提到的月球水中不同的 D/H 比例的测量相矛盾,大约是地球海洋中的两倍(Greenwood 等人,2011)。如果有效,这将表明在形成月球的撞击后不久,对地月系统进行了大量彗星水的输送。因此,地球水将是一个晚期添加,仅由一个或最多几次与错过月球的地球碰撞导致(Robert,2011)。因此,扩大具有准确 D/H 测量的对象样本是新的亚毫米设施的一个高优先级,长期科学目标。

While the case for late delivery of water on Earth is still open, Lunine et al. (2003) showed that the abundance of water derived for early Mars is consistent with the general picture of late delivery by comets and bodies from the early asteroid belt. Unlike Earth, however, the small size of Mars dictates that it acquired its water primarily from very small bodies—asteroidal (sizes up to 100 km)—rather than lunar (1700\sim 1700 km) or larger. The stochastic nature of the accretion process allows this to be one outcome out of many, but the Jupiter Grand Tack scenario provides a specific mechanism for removing larger bodies from the region where Mars formed, thereby resulting in its small size.
尽管地球上水的延迟交付的情况仍然存在争议,但 Lunine 等人(2003 年)表明,早期火星的水丰度与晚期由彗星和来自早期小行星带的天体交付的一般情况一致。然而,与地球不同,火星的小尺寸决定了它主要从非常小的天体(尺寸最多达 100 公里的小行星)而不是月球( 1700similar-toabsent1700\sim 1700 公里)或更大的天体中获取水。聚集过程的随机性使得这可以成为许多结果之一,但木星的大迁移场景提供了一种特定的机制,可以将火星形成区域的较大天体移除,从而导致其小尺寸。

8.2.2 Wet scenario  8.2.2 湿润场景

There are two ways for Earth to get its water locally around 1 AU rather than through delivery from the outer solar system. The first option is that local planetesimals have retained some water at high temperatures through chemisorption onto silicate grains (Drake, 2005; de Leeuw et al., 2010). There is no evidence for such material today in meteorites nor in the past in interstellar and protoplanetary dust (§ 5.1). However, the solid grains were bathed in abundant water vapor during the entire lifetime of gas-rich disks. This mechanism should differentiate between Earth and Venus. Computation and lab studies (Stimpfl et al., 2006) suggest that chemisorption may be just marginally able to supply the inventory needed to explain Earth’s crustal water, with some mantle water, but moving inward to 0.7 AU the efficiency of chemisorption should be significantly lower. Finding that Venus was significantly dryer than the Earth early in the solar system history would argue for the local wet source model, although such a model may also overpredict the amount of water acquired by Mars.
地球获得水的两种方式是在 1 AU 附近本地获取,而不是通过外太阳系的输送。第一种选择是本地的小行星在高温下通过化学吸附到硅酸盐颗粒上保留了一些水。今天的陨石中没有这样的材料的证据,也没有在星际和原行星尘埃中发现过。然而,固体颗粒在富含气体的盘的整个寿命期间都被大量水蒸气浸泡。这种机制应该能区分地球和金星。计算和实验室研究表明,化学吸附可能仅能够边缘地提供解释地球地壳水以及一些地幔水所需的库存,但向内移动到 0.7 AU,化学吸附的效率应该显著降低。发现金星在太阳系早期比地球干燥得多将支持本地湿源模型的论点,尽管这种模型也可能高估了火星获得的水量。

The second wet scenario is that Earth accretes a water rich atmosphere directly from the gas in the inner disk. This would require Earth to have had in the past a massive hydrogen atmosphere (with a molar ratio H2/H2O larger than 1) that experienced a slow hydrodynamical escape (Ikoma and Genda, 2006). One problem with this scenario is that the timescales for terrestrial planet formation are much longer than the lifetime of the gas disk. Also, the D/H ratio would be too low, unless photochemical processes at the disk surface enhance D/H in water and mix it down to the midplane (Thi et al., 2010a) or over long timescales via mass-dependent atmospheric escape (Genda and Ikoma, 2008).
第二种潮湿的情景是地球直接从内部盘的气体中吸收富含水的大气层。这将需要地球在过去拥有一个巨大的氢气大气层(其摩尔比 H 2 /H 2 O 大于 1),并经历了缓慢的流体动力逃逸(Ikoma 和 Genda,2006)。这种情景的一个问题是,地球类行星形成的时间尺度远远长于气体盘的寿命。此外,除非光化学过程在盘面上增强了水的 D/H 比率并将其混合到中平面(Thi 等人,2010a),或者通过质量依赖的大气逃逸(Genda 和 Ikoma,2008)长时间尺度上混合,否则 D/H 比率将太低。

9 EXOPLANETARY ATMOSPHERES
9 外行星大气

H2O is expected to be one of the dominant components in the atmospheres of giant exoplanets, so searches for signatures of water vapor started immediately when the field of transit spectroscopy with Spitzer opened up in 2007. Near-IR spectroscopy with HST and from the ground can be a powerful complement to these data because of the clean spectral features at 1–1.6 μ\mum (Fig. 1). Early detections of water in HD 189733b, HD 209458 and XO1b were not uniformly accepted (see Seager and Deming, 2010; Tinetti et al., 2012, and chapter by Madhusudhan et al. for summaries). Part of the problem may stem from hazes or clouds that can affect the spectra shortward of 1.5 μ\mum and hide water which is visible at longer wavelengths. Even for the strongest cases for detection, it is difficult to retrieve an accurate water abundance profile from these data, because of the low S/NS/N and low spectral resolution (Madhusudhan and Seager, 2010). The advent of a spatial scan mode has improved HST’s ability to detect exoplanetary water with WFC3 (e.g., Deming et al., 2013; Mandell et al., 2013). New ground-based techniques are also yielding improved spectra (Birkby et al., 2013), but dramatic improvements in sensitivity and wavelength range must await JWST. Until that time, only qualitative conclusions can be drawn.
H 2 O 预计将成为巨型外行星大气中的主要组分之一,因此当 2007 年 Spitzer 的过境光谱学领域开放时,对水蒸气迹象的搜索立即开始。HST 和地面的近红外光谱学可以成为这些数据的强大补充,因为在 1-1.6 μ\mu m 处具有清晰的光谱特征(图 1)。在 HD 189733b、HD 209458 和 XO1b 中早期对水的探测并未被一致接受(参见 Seager 和 Deming,2010 年;Tinetti 等人,2012 年,以及 Madhusudhan 等人的章节摘要)。问题的一部分可能源于可能影响 1.5 μ\mu m 以下光谱的雾霾或云层,这些云层会隐藏在较长波长处可见的水。即使对于检测最强烈的情况,由于低 S/NS/N 和低光谱分辨率(Madhusudhan 和 Seager,2010 年),从这些数据中检索准确的水含量剖面也是困难的。空间扫描模式的出现已经改善了 HST 使用 WFC3 检测外行星水的能力(例如,Deming 等人,2013 年;Mandell 等人,2013 年)。 新的地面技术也正在提供改进的光谱(Birkby 等人,2013 年),但灵敏度和波长范围的显著改进必须等待 JWST。在那之前,只能得出定性结论。

Refer to caption
Figure 13: Spectra of the hot Jupiters with C/O ratios varying from 0.54 (solar, red) to 1 (blue) and 3 (green). Note the disappearance of the H2O features at 1–3 μ\mum as the C/O ratio increases. Figure from Madhusudhan et al. (2011b).
图 13:热木星的光谱,其碳氧比从 0.54(太阳,红色)变化到 1(蓝色)和 3(绿色)。请注意,随着碳氧比的增加,H 2 O 特征在 1-3 μ\mu m 处消失。图片来源于 Madhusudhan 等人(2011b)。

9.1 Gas-giant planets  9.1 气态巨行星

Models of hydrogen-dominated atmospheres of giant exoplanets (T7003000T\approx 700-3000 K; warm Neptune to hot Jupiter) solve for the molecular abundances (i.e., ratios of number densities, also called ‘mixing ratios’) under the assumption of thermochemical equilibrium. Non-equilibrium chemistry is important if transport (e.g., vertical mixing) is rapid enough that material moves out of a given temperature and pressure zone on a timescale short compared with that needed to reach equilibrium. The CH4/CO ratio is most affected by such ‘disequilibrium’ chemistry, especially in the upper atmospheres. Since CO locks up some of the oxygen, this can also affect the H2O abundance.
巨大外行星( T70030007003000T\approx 700-3000 K;温暖的海王星到热木星)的氢为主的大气模型解决了分子丰度(即,数密度比,也称为“混合比率”)在热化学平衡假设下的问题。如果传输(例如,垂直混合)足够快,使得物质在短于达到平衡所需时间的时间尺度内移出给定的温度和压力区域,那么非平衡化学就很重要。CH 4 /CO 比率在这种“非平衡”化学中受到最大影响,尤其是在上层大气中。由于 CO 锁定了一些氧气,这也会影响 H 2 O 的丰度。

The two main parameters determining the water abundance profiles are the temperature-pressure combination and the C/O ratio. Giant planets often show a temperature inversion in their atmospheres, with temperature increasing with height into the stratosphere, which affects both the abundances and the spectral appearance. However, for the full range of pressures (
确定水丰度剖面的两个主要参数是温度-压力组合和 C/O 比。巨大行星的大气层通常显示温度反转,随着高度增加,温度逐渐升高进入平流层,这影响了丰度和光谱外观。然而,对于全压力范围(
10210^{2}10510^{-5} bar) and temperatures (700–3000 K), water is present at abundances greater than
bar)和温度(700-3000 K),水的丰度大于
10310^{-3} with respect to H2 for a solar C/O abundance ratio of 0.54, i.e., most of the oxygen is locked up in H2O. In contrast, when the C/O ratio becomes close to unity or higher, the H2O abundance can drop by orders of magnitude especially at lower pressures and higher temperatures, since the very stable CO molecule now locks up the bulk of the oxygen
相对于 H 2 ,对于太阳 C/O 丰度比 0.54,即大部分氧被锁定在 H 2 O 中。相反,当 C/O 比接近或更高时,H 2 O 的丰度可以在较低压力和较高温度下大幅下降,因为非常稳定的 CO 分子现在锁定了大部分氧
(Kuchner and Seager, 2005; Madhusudhan et al., 2011b)
(Kuchner 和 Seager,2005;Madhusudhan 等,2011b)
. CH4 and other hydrocarbons are enhanced as well. As noted previously, observational evidence (not without controversy) for at least one case of such a carbon-rich giant planet, WASP-12b, has been found
. CH 4 和其他碳氢化合物也得到增强。正如先前所指出的,至少有一个这样的碳丰富巨行星 WASP-12b 的观测证据(尽管有争议)已经被发现
(Madhusudhan et al., 2011a)
(Madhusudhan 等,2011a)
. Hot Jupiters like WASP-12b are easier targets for measuring C/O ratios, because their elevated temperature profiles allow both water and carbon-bearing species to be measured without the interference from condensation processes that hamper measurements in Jupiter and Saturn in our own solar system, although cloud-forming species more refractory than water ice can reduce spectral contrast for some temperature ranges. Figure 13 illustrates the changing spectral appearance with C/O ratio for a hot Jupiter (
. 像 WASP-12b 这样的热木星更容易用于测量 C/O 比值,因为它们升高的温度轮廓允许测量水和含碳物质,而不受干扰来自凝结过程的影响,这些过程妨碍了在我们太阳系中的木星和土星中的测量,尽管比水冰更难挥发的云形成物质可能会降低某些温度范围的光谱对比度。图 13 说明了热木星随着 C/O 比值的变化而产生的光谱外观 (
T=2300T=2300 K); for cooler planets (
K);对于较冷的行星 (
<1000<1000 K) the changes are less obvious.
变化不太明显。

What can cause exoplanetary atmospheres to have very different C/O ratios from those found in the interstellar medium or even in their parent star? If giant planets indeed form outside the snow line at low temperatures through accretion of planetesimals, then the composition of the ices is a key ingredient in setting the C/O ratio. The volatiles (but not the rocky cores) vaporize when they enter the envelope of the planet and are mixed in the atmosphere during the homogenization process. As the giant planet migrates through the disk, it encounters different conditions and thus different ice compositions as a function of radius (Mousis et al., 2009; Öberg et al., 2011). Specifically a giant planet formed outside the CO snow line at 20–30 AU around a solar mass star may have a higher C/O ratio than that formed inside the CO (but outside the H2O) snow line.
什么会导致外行星大气具有与星际介质甚至其母星中发现的 C/O 比率非常不同的特征?如果巨大行星确实是在低温下通过行星碎片的聚集在雪线外形成的,那么冰的组成是设定 C/O 比率的关键因素。挥发物质(但不包括岩心)进入行星的包层时蒸发,并在均质化过程中混合在大气中。当巨大行星穿越盘时,它会遇到不同的条件,因此不同的冰组成会随着半径的变化而变化(Mousis 等,2009 年;Öberg 等,2011 年)。具体来说,围绕太阳质量星体在 20-30 AU 处形成的超级行星可能比在 CO(但在 H 2 O)雪线内部形成的行星具有更高的 C/O 比率。

The exact composition of these planetesimals depends on the adopted model. The solar system community traditionally employs a model of the solar nebula disk, which starts hot and then cools off with time allowing various species to re-freeze. In this model, the abundance ratios of the ices are set by their thermodynamic properties (‘condensation sequence’) and the relative elemental abundances. For solar abundances, water is trapped in various clathrate hydrates like NH3-H2O and H2S-5.75 H2O between 5 and 20 AU. In a carbon-rich case, no water ice is formed and all the oxygen is CO, CO2 and CH3OH ice (Madhusudhan et al., 2011b). The interstellar astrochemistry community, on the other hand, treats the entire disk with non-equilibrium chemistry, both with height and radius. Also, the heritage of gas and ice from the protostellar stage into the disk can be considered (Aikawa and Herbst, 1999; Visser et al., 2009; Aikawa et al., 2012; Hincelin et al., 2013, see Fig. 5). The non-equilibrium models show that indeed the C/O ratio in ices can vary depending on location and differ by factors of 2–3 from the overall (stellar) abundances (Öberg et al., 2011).
这些小行星的确切组成取决于所采用的模型。太阳系社区传统上采用太阳星云盘模型,该模型一开始很热,然后随着时间的推移冷却,使得各种物种可以重新结晶。在这个模型中,冰的丰度比由它们的热力学性质(“凝结序列”)和相对元素丰度确定。对于太阳丰度,水被困在各种类沸石水合物中,如 NH 3 -H 2 O 和 H 2 S-5.75 H 2 O 在 5 到 20 AU 之间。在富含碳的情况下,不形成水冰,所有氧都是 CO,CO 2 和 CH 3 OH 冰(Madhusudhan 等人,2011b)。另一方面,星际天体化学社区对整个盘进行非平衡化学处理,无论是高度还是半径。此外,可以考虑从原恒星阶段到盘中的气体和冰的遗传(Aikawa 和 Herbst,1999;Visser 等人,2009;Aikawa 等人,2012;Hincelin 等人,2013,见图 5)。 非平衡模型表明,冰中的 C/O 比确实会因位置而异,并且与整体(恒星)丰度相差 2-3 倍(Öberg 等人,2011)。

9.2 Rocky planets, super-Earths
9.2 岩石行星,超类地

The composition of Earth-like and super-Earth planets (up to 10 M\EarthM_{\Earth}) is determined by similar thermodynamic arguments, with the difference that most of the material stays in solid form and no hydrogen-rich atmosphere is attracted. Well outside the snow line, at large distances from the parent star, the planets are built up largely from planetesimals that are half rock and half ice. If these planets move inward, the water can become liquid, resulting in ‘ocean planets’ or ‘waterworlds’, in which the entire surface of the planet is covered with water (Kuchner, 2003; Léger et al., 2004). The oceans on such planets could be hundreds of kilometers deep and their atmospheres are likely thicker and warmer than on Earth because of the greenhouse effect of water vapor. Based on their relatively low bulk densities (from the mass-radius relation), the super-Earths GJ 1214b (Charbonneau et al., 2009) and Kepler 22b (Borucki et al., 2012) are candidate ocean planets.
类地和超类地行星(最多可达 10 个 M\EarthsubscriptM_{\Earth} )的组成由类似的热力学论证确定,不同之处在于大部分材料保持固态形式,不会吸引富氢大气。在远离母恒星的大距离处的雪线外,这些行星主要由半岩石半冰的小行星堆积而成。如果这些行星向内运动,水可能变成液态,形成“海洋行星”或“水世界”,其中行星的整个表面都被水覆盖(Kuchner, 2003; Léger 等,2004)。这些行星上的海洋可能有数百公里深,它们的大气可能比地球上的大气更厚更温暖,这是由于水蒸气的温室效应。根据它们相对较低的总体密度(从质量-半径关系),超类地 GJ 1214b(Charbonneau 等,2009)和 Kepler 22b(Borucki 等,2012)是海洋行星的候选者。

Further examination of terrestrial planet models shows that there could be two types of water worlds: those which are true water rich in their bulk composition and those which are mostly rocky but have a significant fraction of their surface covered with water (Earth is in the latter category) (Kaltenegger et al., 2013). Kepler 62e and f are prototypes of water-rich planets within the habitable zone, a category of planets that does not exist in our own solar system (Borucki et al., 2013). Computing the atmospheric composition of terrestrial exoplanets is significantly more complex than that of giant exoplanets and requires consideration of many additional processes, including even plate tectonics (Meadows and Seager, 2011; Fortney et al., 2013).
对地球类行星模型的进一步研究表明,可能存在两种类型的水世界:一种是在其主要成分中富含水的真正水世界,另一种是主要为岩石但表面有大量水覆盖的水世界(地球属于后者)(Kaltenegger 等人,2013 年)。Kepler 62e 和 f 是宜居带内富含水的行星的原型,这是我们太阳系中不存在的一类行星(Borucki 等人,2013 年)。计算类地外行星的大气组成比计算巨行星的复杂得多,需要考虑许多额外的过程,甚至包括板块构造(Meadows 和 Seager,2011 年;Fortney 等人,2013 年)。

As for the terrestrial planets in our own solar system, the presence of one or more giant planet can strongly affect the amount of water on exo-(super)-Earths (Fig. 12). A wide variety of dynamical and population synthesis models on possible outcomes have been explored (e.g. Raymond et al., 2006; Mandell et al., 2007; Bond et al., 2010; Ida and Lin, 2004, 2010; Alibert et al., 2011; Mordasini et al., 2009).
至于我们太阳系中的地球类行星,一个或多个巨行星的存在可能会对外星(超级)地球上的水量产生强烈影响(见图 12)。已经探讨了各种可能结果的动力学和群体综合模型(例如 Raymond 等人,2006 年;Mandell 等人,2007 年;Bond 等人,2010 年;Ida 和 Lin,2004 年,2010 年;Alibert 等人,2011 年;Mordasini 等人,2009 年)。

Water on a terrestrial planet in the habitable zone may be cycled many times between the liquid oceans and the atmosphere through evaporation and rain-out. However, only a very small fraction of water molecules are destroyed or formed over the lifetime of a planet like Earth. The total Earth hydrogen loss is estimated to be 3 kg s-1. Even if all the hydrogen comes from photodissociation of water, the water loss would be 27 kg s-1. Given the total mass of the hydrosphere of 1.5×10211.5\times 10^{21} kg, it would take 1.8×10121.8\times 10^{12} yr to deplete the water reservoir. The vast majority of the water bonds present today were, therefore, formed by the chemistry that led to the bulk of water in interstellar clouds and protoplanetary disks.
在宜居带内的地球类行星上,水可能在液态海洋和大气层之间多次循环,通过蒸发和降雨。然而,在类似地球的行星寿命期间,只有极小部分水分子被破坏或形成。据估计,地球总氢损失量为 3 千克每秒。即使所有氢来自水的光解,水的损失量也将为 27 千克每秒。考虑到水圈的总质量为 1.5×10211.5superscript10211.5\times 10^{21} 千克,将需要 1.8×10121.8superscript10121.8\times 10^{12} 年才能耗尽水库。因此,今天存在的绝大多数水键是由导致星际云和原行星盘中大部分水形成的化学反应形成的。

10 WATER TRAIL FROM CLOUDS TO PLANETS
从云到行星的 10 水迹

Here we summarize the key points and list some outstanding questions.
在这里,我们总结了关键要点并列出了一些未解之谜。

  • Water is formed on ice in dense molecular clouds. Some water is also formed in hot gas in shocks associated with protostars, but that water is largely lost to space.
    在密集的分子云中,水在冰上形成。一些水也在与原恒星伴随的冲击热气体中形成,但这些水大部分都流失到太空中。

  • Water stays mostly as ice during protostellar collapse and infall. Only a small fraction of the gas in the inner envelope is in a ’hot core’ where the water vapor abundance is high due to ice sublimation.
    水在原恒星坍缩和坠落过程中主要以冰的形式存在。只有内层包层中的一小部分气体处于“热核”状态,其中水蒸气的丰度较高,这是由于冰升华造成的。

  • Water enters the disk mostly as ice at large radii and is less affected by the accretion shock than previously thought.
    水主要以冰的形式进入盘中,在较大半径处受到的加积冲击影响较小,这比以前认为的要少。

  • Water vapor is found in three different reservoirs in protoplanetary disks: the inner gaseous reservoir, the outer icy belt and the hot surface layers. The latter two reservoirs have now been observed with Spitzer and Herschel and quantified.
    在原行星盘中,水蒸气存在于三个不同的储存库中:内部气态储存库、外部冰带和热表面层。后两个储存库已经被 Spitzer 和 Herschel 观测到并进行了量化。

  • Models suggest that the water that ends up in the planet and comet forming zones of disks is partly pristine ice and partly processed ice, i.e., ice that has at least once sublimated and recondensed when the material comes close to the young star.
    模型表明,在行星和彗星形成区域的盘中最终存在的水部分是原始冰和经过加工的冰,即至少曾经在物质靠近年轻恒星时升华和重新凝结的冰。

  • Water ice promotes grain growth to larger sizes; water-coated grains grow rapidly to planetesimal sizes in dust traps that have now been observed.
    冰水促进颗粒生长到更大的尺寸;涂有水的颗粒在尘埃陷阱中迅速增长到行星大小,这已经被观察到。

  • Planetesimals inside the disk’s snow line are expected to be dry, as found in our solar system. However, gaseous water can have very high abundances in the inner AU in the gas-rich phase. Some grains may have chemically bound water to silicates at higher temperatures, but there is no evidence of hydrated silicates in meteorites nor in interstellar grains.
    盘内雪线内的小行星预计是干燥的,就像在我们的太阳系中发现的那样。然而,在富含气体的相位中,气态水可能具有非常高的丰度。一些颗粒可能在较高温度下将水化学地结合到硅酸盐中,但在陨石中或星际颗粒中没有水合硅酸盐的证据。

  • Dynamics affects what type of planetesimals are available for planet formation in a certain location. The presence and migration of giant planets can cause scattering of water-rich planetesimals from the outer disk into the inner dry zone.
    动力学影响了某一位置可用于行星形成的小行星类型。巨行星的存在和迁移可能导致富含水的小行星从外盘散射到内干区。

  • Several comets have now been found with D/H ratios in water consistent with that of Earth’s oceans. This helps to constrain models for the origin of water on the terrestrial planets, but does not yet give an unambiguous answer.
    现在已经发现几颗彗星的 D/H 比与地球海洋一致。这有助于限制关于地球类行星水源起源的模型,但尚未给出明确答案。

  • Both dry and wet formation scenarios for Earth are still open, although most arguments favor accretion of water-rich planetesimals from the asteroid belt during the late stages of terrestrial planet formation. Delivery of water through bombardment to planetary surfaces in a ‘late veneer’ can contribute as well but may not have been dominant for Earth, in contrast with Mars.
    地球的干湿形成场景仍然存在,尽管大多数论点支持在地球类行星形成的后期阶段从小行星带吸积富含水的小行星。通过‘晚期外衣’对行星表面进行轰击传递水分也可能起到一定作用,但与火星相比,可能不是地球的主要机制。

  • Jupiter may be poor in oxygen and water. If confirmed by the Juno mission to Jupiter, this may indicate a changing C/O ratio with disk radius. Evidence for this scenario comes from the detection of at least one carbon-rich exoplanet.
    木星可能缺乏氧气和水。如果由朱诺任务对木星的确认,这可能表明随着盘半径的变化,C/O 比例也在变化。支持这种情况的证据来自至少一个富含碳的外行星的探测。

  • All of the processes and key parameters identified here should also hold for exo-planetary systems.
    这里确定的所有过程和关键参数也适用于外行星系统。

There are a number of open questions that remain. A critical phase is the feeding of material from the collapsing core onto the disk, and the evolution of the young disk during the bulk of the phase of star formation. At present this phase has little observational constraints. In addition numerous models posit that viscous evolution and the relative movement of the dust to the gas can have impact on the overall water vapor evolution; this needs an observational basis. Astronomical observations have been confined to water vapor and ice emission from the disk surface but with the midplane hidden from view. What does this surface reservoir tells us about forming planets, both terrestrial and giant? Is there a way to detect the midplane, perhaps using HCO+ (which is destroyed by water) as a probe of the snowline? However, this needs a source of ionization in the densest parts of the disk which may or may not be present (Cleeves et al., 2013).
仍然存在许多未解之谜。关键阶段是从坍缩的核心向盘上的物质供给,以及在恒星形成的大部分阶段期间年轻盘的演化。目前这个阶段几乎没有观测约束。此外,许多模型认为粘性演化和尘埃相对于气体的运动可能对整体水蒸气演化产生影响;这需要有观测基础。天文观测一直局限于盘面的水蒸气和冰的发射,但中面隐藏在视野之外。这个表面储库告诉我们关于形成行星的信息,无论是类地行星还是巨行星?是否有一种方法可以探测中面,也许可以使用 HCO(被水破坏)作为雪线的探针?然而,这需要在盘的最密集部分有电离源,这可能存在也可能不存在(Cleeves 等人,2013)。

The exploration of deuterium enrichments continues to hold promise both in the solar system, with need for more information on the D/H values for water in the outer asteroids and (main belt) comets. It is also clear that we know less about the bulk elemental abundance of the solar system’s largest reservoir of planetary material (i.e., Jupiter and Saturn) than one might have assumed. This must have implications for studies of extra-solar systems. In terms of the origin of the Earth’s oceans matching the full range of geochemical constraints remains difficult.
氘富集的探索继续展现出潜力,尤其是在太阳系中,需要更多关于外部小行星和(主带)彗星水的 D/H 值的信息。很明显,我们对太阳系中最大的行星物质库(即木星和土星)的总体元素丰度了解较少,这对外太阳系的研究可能有影响。就地球海洋的起源而言,要使其与地球化学约束的全部范围相匹配仍然很困难。

Despite the uncertainties, there is a better understanding of the water trail from the clouds to planets. In this light, it is fascinating to consider that most of the water molecules in Earth’s oceans and in our bodies may have been formed 4.6 billion years ago in the cloud out of which our solar system formed.
尽管存在不确定性,对从云到行星的水迹象有了更好的理解。在这种光下,考虑到地球海洋和我们身体中大部分水分子可能在 46 亿年前形成于云中,而我们的太阳系也是在这些云中形成的,这是非常迷人的。

Acknowledgments. The authors thank many colleagues for collaborations and input, and various funding agencies for support, including NASA/JPL/Caltech. Figures by Magnus Persson, Ruud Visser, Lars Kristensen and Davide Fedele are much appreciated.
致谢。作者感谢许多同事的合作和贡献,以及包括 NASA/JPL/Caltech 在内的各种资助机构的支持。Magnus Persson、Ruud Visser、Lars Kristensen 和 Davide Fedele 制作的图表受到了很高的赞赏。


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