Abstract– The rock record from the early solar system indicates high-temperature thermal processing sufficient to melt refractory oxides and silicates. The astrophysical context for the formation and evolution of our solar system, from a molecular cloud to a “clean” planetary system, is difficult to constrain tightly because of the large scales and lack of resolution of astronomical observations. Protostellar jets and winds, commonly associated with forming stars, are likely to play a role in heating and redistribution of the processed material in the solar system. We have recently proposed that disk-winds can cause melting of small inclusions to distances out to several AU. Particularly energetic outbursts, such as the FU-Orionis and EXor events, occur over relatively short time scales (approximately 100 and 1 yr, respectively), and are probably events related to formation of the refractory solids present in primitive meteorites.
The formation of planets around stars in the Milky Way galaxy now appears to be the rule rather than the exception (Cassan et al. 2012). This suggests a common mechanism for the formation of solar systems like our own. However, there are still many unresolved issues concerning the processes that led to the formation of planetary systems. Observing protoplanetary disks is difficult because it is hard to achieve the required resolution to “see” into the inner disk that encircles the young star (i.e., inside the equivalent of several astronomical units [AU] from the protostar) and because of scattering of the direct light from the central star. Furthermore, astronomical observations often lack a sufficiently precise temporal constraint, and the age of the protostellar object can only be inferred from its properties relative to a standard model (e.g., Stahler 1983; D’Antona and Mazzitelli 1994).
On the other hand, the meteoritic record of our own solar system provides a detailed view of some of the constituents of the protoplanetary disk associated with the young Sun. These products include meteorites which contain refractory materials, evidence suggestive of high-temperature processing. However, the meteoritic evidence lacks the physical context in which this processing took place. It is unclear where the different components of meteorites were processed, and we only have a first-order knowledge of when this happened within the astronomical context.
It is evident that one of the key issues preventing progress in this field is that astronomy has too little spatial definition, and meteoritics has too much. In astronomy, we cannot resolve the activities in the disk at distances less than tens of AU; in meteoritics, we are looking at objects that grow to a few centimeters, and are effectively unconstrained in their location of formation.
Even so, the joint consideration of astrophysical and meteoritical constraints over the past decade has begun to show progress in deciphering the processes that led to the formation of our solar system. This is a fruitful approach, expected to drive substantial advances into the future.
In this article, we begin by examining the star-formation process from an astronomical perspective. We then examine features from the meteoritic record that show evidence for the high-temperature processing of magnesium-silicate droplets (chondrules), and for the cold accretion of chondritic meteorites. We also consider the role of protostellar jets, commonly associated with young stars, in the processing and transportation of this material. We argue that extended disk-winds may provide a mechanism for the formation of chondrules in close association with the cold material of the protostellar disk.
Star Formation and Astrophysical Chronology
New stars are born out of the gas and dust that collects in dense “cores” of interstellar clouds (e.g., see the reviews by Shu et al. 1987; Cameron 1995; André et al. 2000; Larson 2003; McKee and Ostriker 2007). These clouds coalesce in the interstellar medium, forming large structures that can be observed in the spiral arms of galaxies. These structures may contain several giant molecular clouds of up to a million solar masses in size and exhibiting filamentary or clumpy substructures over a wide range of spatial scales (e.g., see Lada et al. 2007).
Core formation is thought to result from the gradual loss of magnetic support against the cloud self-gravity, in turn due to the slow outward diffusion of magnetic fields through the cloud material (see e.g., Mouschovias et al. 2006). The density profile of the core develops a central peak, which eventually becomes gravitationally unstable and initiates an “inside-out” dynamical collapse. Alternatively, supersonic turbulent motions may play a role in the onset of cloud collapse (e.g., Mac Low and Klessen 2004; Ballesteros-Paredes et al. 2007). Regardless of the specific triggering mechanism, during the ensuing phase of dynamical collapse, the specific angular momentum of cloud material is approximately conserved as it flows toward the center of the core. This results in an increase in its rotational velocity and in the associated centrifugal force, until this force is strong enough to balance the inward pull of gravity. The infalling material then settles into a rotationally supported disk, on scales of about 100 AU around the forming star (see Fig. 1). Disk matter subsequently accretes onto the protostar, provided that it can transport away its excess angular momentum.
The main collapse phase of the star-formation process lasts around a few hundred thousand years, during which about half of the mass of the star is assembled (e.g., Shu et al. 1987). During this stage, the infall rate from the cloud onto the disk typically exceeds the accretion rate onto the protostar, as they are regulated by different physical processes (Kenyon et al. 1990; Calvet et al. 2000). This results in the accumulation of mass in the disk, which could eventually become gravitationally unstable (Larson 1984). An episode of enhanced accretion would follow in this case, until the disk dumps enough mass to the center to regain gravitational stability. These episodes of high accretion rate have been associated with the FU-Orionis outbursts exhibited by forming low-mass stars (Kenyon et al. 1990).
Ultimately, disk masses are much smaller than that of the central object, and their lifetime is typically in the range of 1–10 Myr (Calvet et al. 2000). It is also worth noting that these flattened disks of material are also associated with other astrophysical systems undergoing accretion, such as the disks surrounding the cores of active galaxies and galactic black holes.
According to the “core accretion” model of planet formation (Pollack et al. 1996), planetary bodies begin to form by aggregation of solid particles suspended in the disk. These dust grains aggregate rapidly via collisions and sink to the disk’s equatorial plane (i.e., the “midplane”), where the resulting thin layer of dust becomes unstable to fragmentation (i.e., formation of clumps) under its own gravity. Provided that their growth proceeds faster than their collisional destruction, these aggregates become bigger, and eventually begin to attract other aggregates by mutual gravitational attraction (e.g., the formation of the rubble-pile asteroid Itokawa; Fujiwara 2006). However, how growth beyond around a meter is achieved is not well understood, as for 1 m-sized particles destructive processes are believed to occur more efficiently than those leading to growth (e.g., Dominik et al.  and references therein), and because these bodies experience strong radial drift and are expected to move rapidly inwards (e.g., Nagasawa et al. 2007). Self-evidently, this process results in kilometer-sized rocky bodies which grow further through collisions, until a few dominant planetary cores emerge. Some of these cores can grow quickly to a size where their gravitational force is sufficient to capture the neighboring gas, yielding planets of Jupiter size and beyond.
The Solar Nebula
A commonly used representation of the physical conditions of the early solar system in the astrophysical community is the minimum-mass solar nebula model (MMSN; Hayashi 1981). In this model, the radial “surface density” profile of the nebular gas (i.e., the vertically integrated density as a function of radius) is arrived at by supplying enough hydrogen and helium to the mass of the current planetary bodies of the solar system, until recovering the standard elemental abundances. The resulting mass is then spread in a disk-like shape. The ambient temperature (assumed to be independent of height) is obtained by assuming that the gas is in thermal balance with the solar radiation field. The midplane density can be estimated under the assumption that the disk is in vertical hydrostatic equilibrium (i.e., the vertical component of the solar gravitational force balances the vertical pressure gradient in the disk). The obtained surface density, midplane density, and temperature radial profiles of the solar nebula disk are given, respectively, by the following expressions, where RAU denotes the radius in AU (see also Fig. 2):
According to the minimum-mass solar nebula model, the surface density of the protosolar disk at 1, 5, and 10 AU is 1700, 150, and 50 g cm−2, respectively. Note, however, that astronomical observations suggest that the surface density in the inner regions of protostellar disks may be lower than these values, and that this quantity may decline more gradually with radius (Kitamura et al. 2002; Andrews and Williams 2007). On the other hand, protostellar disks may have surface densities of up to about 30 times these values—between 1 and 100 AU—without being gravitationally unstable (Durisen et al. 2007). Importantly, it has been pointed out that calculated disk masses may be systematically underestimated (e.g., see Hartmann et al. 2006; Williams and Cieza 2011; Jones et al. 2012). This is because disk masses are typically estimated from measurements of the emission from dust grains, together with assumptions about the gas-to-dust mass ratio and the relative abundance of smaller grains (which dominate the radiative emission). Consequently, processes such as grain growth and settling (which reduce the dust opacity) may lead to miscalculated (lower) disk masses.
The MMSN temperature profile given above indicates that temperatures are always below 280 K beyond 1 AU, and increase to about 440 K at the orbital location of Mercury (approximately 0.4 AU). More generally, however, the temperature structure of protostellar disks is the result of a number of complex heating and cooling processes acting in the disk (e.g., see the reviews by Dullemond et al.  and Najita et al. ). Both passive disks, heated only by irradiation from the central object (e.g., Dullemond et al. 2002), and actively accreting viscous disks, where viscous dissipation is an additional heating source (e.g., D’Alessio et al. 2006) have been investigated. In passive-disk models, the disk temperature is found to increase from the midplane (where it is in the range of 10–60 K for radii decreasing from 300 to 1 AU) to the surface layers (where it plateaus at a value of 30–200 K for the same radial extension; Dullemond et al. 2002). When the disk is also heated by viscous dissipation, the temperature in the disk interior may initially drop away from the midplane, reaching a minimum at an intermediate height, and then rising again toward the surface (D’Alessio et al. 2006). As these authors point out, the actual temperature profile depends on the size distribution and degree of settling of the grains embedded in the gas. For these types of models, the midplane and surface temperatures are in the range of 20–250 K and 50–400 K, respectively, at radial scales of 1–100 AU.
The above studies adopt the common assumption that gas and dust are thermally coupled and, therefore, have the same temperature. This is a good approximation in the disk interior, where the mutual collisions between these species are frequent, but it breaks down toward the surface, where the density is too low to enforce this coupling. In the surface layers therefore the gas is expected to become significantly hotter than the dust. Glassgold et al. (2004) find that in an active, viscous disk, at 1 AU from the center, the gas temperature reached approximately 5000 K in the disk atmosphere (at a vertical column density less than 1021 g cm−2), whereas the dust temperature remained below 400 K. Exploring the temperature structure of a wind-driving disk, which is heated magnetically, and studying the impact of magnetic activity on the dynamical and thermal processing of dust particles embedded in it, is a distinctly fertile area of research.
Protostellar Jets and Disk-Winds
A key feature of young stars is their association with supersonic, highly collimated winds that propagate along the polar axis of the protostar-disk system. These outflows are detected in protostellar sources over a wide range of masses, from the very early stages of formation (sources associated with infalling envelopes whose mass exceeds that of the growing star) all the way to the so-called “Classical T-Tauri” stars, whose envelopes have already dispersed (e.g., see Arce et al. 2007; Bally et al. 2007). These jets are thought to play a key role in regulating the star-formation process by removing the excess angular momentum of disk material and enabling matter to flow toward the center (e.g., see the reviews by Königl and Pudritz 2000; Pudritz et al. 2007; Königl and Salmeron 2011). The detected correlation between accretion and outflow activity in these objects, with the inferred outflow rates being approximately 0.01–0.1 times the corresponding accretion rates, supports this scenario (Hartigan et al. 1995; Cabrit 2007a). Protostellar winds also provide an important feedback mechanism between the forming star and the surrounding medium, to which they return mass and energy (e.g., Arce et al. 2007).
The outflows accelerated from young stars comparable in size to the Sun typically exhibit terminal velocities along the jet axis in the range of 150–400 km s−1, with lower velocity material located in surrounding concentric rings (this is often referred to as the “onion-like” structure of protostellar jets). This slower material forms a low-velocity sheath with typical velocities of the order of 10–100 km s−1 (Cabrit [2007a] and references therein). Another striking feature of protostellar jets is their high level of collimation, with typical opening angles of the order of only 3–5° even at close distances to the source (i.e., about 50 AU above the disk; Cabrit 2007b). Finally, although the radial extent of the jet-launching region is not fully established, observations using the Hubble Space Telescope Imaging Spectrograph (Coffey et al. 2004, 2007) suggest that the higher velocity component of the outflow from some protostellar sources originates from a radial location only approximately 0.2–0.5 AU from the axis, whereas the low-velocity component may be launched from a radius of up to approximately 2 AU from the protostar.
Protostellar winds may have several origins including the star itself, the stellar magnetosphere, and the disk (e.g., Ferreira et al. 2006) (see Fig. 1). Here, we focus on disk-winds because they may originate from a radial distance encompassing the inferred chondrite-formation region. Disk-winds are thought to accelerate from the disk surfaces by the action of magnetic fields, via the mechanism first proposed by Blandford and Payne (1982; see Fig. 1). According to this picture for the origin of protostellar outflows, the disk is threaded by large-scale, ordered magnetic field lines that bend radially outwards. If these field lines are sufficiently inclined, such that the angle between the line and the disk surface is less than 60°, and the fluid couples to the magnetic field, material near the disk surface can be accelerated centrifugally along the field lines (through the “bead-on-a-wire” effect). This constitutes the base of a large-scale outflow through which material can potentially escape from the system. The existence of magnetic fields with the required topology for this mechanism to operate is reasonably expected both from observational evidence of star-forming regions (Girart et al. 2006) as well as from models of disk formation via the collapse of molecular-cloud cores (Mouschovias 1991; Krasnopolsky and Königl 2002; Kunz and Mouschovias 2009).
Formation of the Solar System
Direct evidence of the formation events of our own solar system is preserved in the meteorite record. Studies of meteorite samples therefore complement analytical and numerical models of the early solar system, as well as observational studies of current star formation. Chondrite meteorites, in particular, are regarded as being the building blocks of the solar system, the material out of which the planets formed. Furthermore, they are inferred to have formed through the clumping process required for planetary formation.
Chondrites have three main components (e.g., Alexander et al.  and references therein): chondrules; calcium-aluminum-rich inclusions (CAIs); and matrix (see Fig. 3). Chondrules are millimeter-sized spherules made of magnesium-silicate minerals. Their properties suggest that they melted at temperatures of order 1800 K (Petaev and Wood 2005; Scott and Krot 2005a). Key properties of CAIs (e.g., their mineralogy, as well as their chemical and isotopic fractionations) suggest that they formed at similar or even higher temperatures (Rubin 2000; Richter et al. 2002; Grossman 2010). However, both of these high-temperature materials are set in a mixture of fine-grained silicates, oxides, organic, and presolar grains (i.e., the matrix) that, as a whole, appear not to have experienced temperatures above approximately 400 K. From a meteoritics perspective therefore the solar nebula seems to be composed of spatially, and/or temporally distinct, thermal regions.
It is worth reiterating here that, according to the overall astrophysical picture of the planet-forming process discussed above, the planets result mainly from aggregation of solid particles present in the nebula. During this process therefore they are not expected to experience significant thermal processing. The inferred ambient temperature of the solar nebula (about 300 K at 1 AU) is consistent with the picture of a relatively cool planet formation environment.
However, high-temperature material is fundamental to the formation of chondrites. Chondrules are not just a minor component of chondrites; they can constitute a large proportion (up to 80%) of these meteorites (Connolly 2005). It follows that a more complete model of the evolution of the solar nebula would need to incorporate a suitable mechanism for the high-temperature processing of refractory materials. Meteoriticists have commonly viewed the solar nebula as a hot environment, as suggested by the mineralogy and melting temperatures of CAIs and chondrules. However, this view has changed as it is at odds with isotopic evidence for a poorly mixed nebula (Liu et al. 2009; Larsen et al. 2011; Simon et al. 2011), and the presence of presolar grains and other low-temperature phases carrying carbon and nitrogen isotope anomalies in carbonaceous material (Anders and Zinner 1993; Nittler 2003). In a high-temperature nebula, these isotope anomalies would be quickly wiped out.
One way to reconcile these divergent pieces of evidence is through transient and localized high-temperature events in an early active solar nebula. Indeed, astronomical observations provide compelling evidence for the occurrence of high-energy processes in young stars, such as the frequent and strong flaring activity in the X-ray emission of protostellar objects (Glassgold et al. 2005), and protostellar jets (e.g., Bally et al. 2007).
Chondrites are the most common type of meteorite (nearly 90% of “falls”; i.e., meteorites that were observed before they reached the ground, and were subsequently collected). They have been divided into three major classes—ordinary, carbonaceous, and enstatite—on the basis of mineralogical, chemical, and oxygen isotope systematics (Krot et al. 2007). Ordinary chondrites constitute the most common type (approximately 90%), and are further subdivided into H, L, and LL chondrites on the basis of their Fe content (respectively: high total iron, low total iron, and low total iron with low metal). Carbonaceous chondrites make up <5% of chondrite meteorites, but are notable in commonly having a large matrix component and relatively large CAI contribution. Finally, enstatite chondrites (about 2%) are notable by having experienced highly reducing conditions, with nearly all Fe present as metal.
The Hayabusa sample return mission to asteroid Itokawa shows that L/LL chondrites are related to the S-class asteroids (Nakamura et al. 2011). This is the second most common class of asteroids, to which 17% of asteroids in the Main Asteroid Belt belong. The dominant C-class asteroids are probably related to carbonaceous chondrites, but confirmation of this relationship might have to wait until the return of the proposed Hayabusa II mission.
Chondrules are subspherical inclusions predominantly composed of olivine, with orthopyroxene, clinopyroxene, and anorthite; minor amounts of troilite (FeS) and Fe, Ni-metal (Connolly 2005). Their textures suggest that these inclusions were once molten, and that these droplets solidified quickly, thereby preserving glass. Chondrules commonly range in size from around 100 μm to 1 mm, although larger chondrules—up to about 10 mm—are also found. The modal abundance of chondrules can be high (up to 80%; Connolly 2005).
Calcium-aluminum-rich inclusions occur predomi-nantly in carbonaceous chondrites. They are mostly composed of minerals such as melilite, Ti-rich pyroxene, spinel, and anorthite in CV3 chondrites; with hibonite, corundum, and perovskite being common in CAIs from CM2 chondrites. Their textures range from “fluffy,” which are inferred to be condensates, to compact and subspherular, which appear to have crystallized from melts. The CV3 CAIs can range in size up to 20 mm, while the CM2 refractory inclusions are seldom larger than 500 μm. CAI abundances are around 3% in CV chondrites and 1% or less in the other carbonaceous chondrites (Hezel et al. 2008).
Chondrules and CAIs are set in a dark, fine-grained matrix. Most chondrite matrices have been seriously affected by parent body alteration, forming Fe-Mg-rich clays, serpentines, and other layered silicate minerals (Buseck and Hua 1993). In very few cases, chondrites have not been affected by this alteration and in these cases the matrices are composed of crystalline Mg-rich silicates, amorphous silicate, presolar grains, refractory solar nebula grains, and rock fragments (Scott and Krot 2005b). Chondrite matrix contains the volatile-rich material which is complementary to the coarse-grained refractory material. The matrix is therefore a diverse assemblage of high-temperature and low-temperature material. These constituents probably formed and were processed in many locations in the solar system (and outside the solar system in the case of the presolar grains).
The proportions of chondrules and matrix can vary quite widely among the different types of chondrites. However, despite the variability in proportions, and the distinctly different chemistry of chondrules and matrix, the chemical composition of each chondrite group is quite constant, suggesting that these components (chondrules and matrix) have a complementary composition, or at least some degree of complementarity in the bulk components (Bland et al. 2005; Zanda et al. 2006; Hezel and Palme 2008; Krot et al. 2009; Desch et al. 2010). This, in turn, could suggest that the components in bulk represent different locations in the nebula, or that there is a close connection between the formation of chondrules at high temperature, and the preservation of the complementary cold matrix in a particular region. The low abundances of CAIs in chondrites make it difficult to assess whether a similar complementary relationship between CAIs and matrix (for refractory elements) could also hold.
Additional evidence for the association of high- and low-temperature material in the early solar system has been inferred from comet 81P/Wild 2 samples collected and brought back to Earth by the Stardust spacecraft. The samples contained crystalline silicates and refractory material similar to chondrules and CAIs (e.g., see Brownlee et al. 2006; McKeegan et al. 2006; Zolensky et al. 2006). Comet Wild 2 originates in the Kuiper Belt (semi-major axis approximately 40 AU). It was therefore expected that the source region of the comet would be dominated by low-temperature material, possibly of molecular-cloud origin, and including interstellar dust (hence the name Stardust for the mission). The recovery of refractory objects was a major surprise. It is unclear whether the crystals found in the samples from Wild 2 formed in the outer protosolar disk (Desch et al. 2005) or in its inner regions, with subsequent transportation to the outer disk (e.g., Ciesla 2011).
Interestingly, crystalline silicates have been detected in the outer regions of protoplanetary disks (Van Boekel et al. 2004; Juhász et al. 2010; see also the review by Henning and Meeus 2011) and in the envelopes that surround young stars in earlier stages of evolution (Poteet et al. 2011). The origin of these crystalline grains is unclear, but because this material seems to be largely absent in the interstellar medium (e.g., Kemper et al. 2004), they are thought to have formed via thermal annealing or evaporation and recondensation of amorphous precursors in the disks and/or envelopes associated with the forming stars. Formation of crystalline silicates has also been reported in recent observations of EX Lupi (Ábrahám et al. 2009), as well as their radial transport via a wind-driven process (Juhász et al. 2012). This protostar is the type-example of a class of accreting young stellar objects characterized by episodic eruptions thought to be driven by an enhancement in the accretion rate onto the central object (e.g., Herbig 2007; Aspin 2011). The observational evidence has been interpreted as showing “in situ” formation of crystalline silicates during the outburst phase (Ábrahám et al. 2009), presumably as a result of the associated increase in the gas temperature.
CAIs and the Early Solar System
Despite their relatively low abundances, CAIs hold a special place in cosmochemistry. They have isotopic compositions that differ from terrestrial (i.e., they appear isotopically anomalous), show trace element abundance patterns that require extremely high temperatures to produce volatility fractionation in even the most refractory elements, and are the oldest objects in the solar system at approximately 4567 Ma.
1 Mass-dependent isotope fractionation (where the relative abundance of an isotope is related to the mass of the isotope), largely caused by physico-chemical processes such as evaporation/condensation of elements like Mg and Si,
2 Radiogenic anomalies related to short-lived nuclides (e.g., 26Mg excesses due to 26Al decay), and
3 Nuclear anomalies, most likely related to distinct nucleosynthetic inputs (e.g., 48Ca, 50Ti).
Calcium-aluminum-rich inclusions are often regarded as the highest temperature objects in the solar system, yet they condense and melt at temperatures similar to chondrules. What sets them apart is their more refractory bulk composition, enriched in Ca, Al, and Ti, forming minerals such as melilite (gehlenite [Ca2Al2SiO7]—akermanite [Ca2MgSi2O7] solid solution), anorthite (CaAl2Si2O8), and perovskite (CaTiO3); as well as their enrichment in refractory trace elements (including the rare earth elements [REE]).
The REE patterns in CAIs are enriched to relatively high levels (around 20 times the value of CI chondrites in CV3 CAIs, and up to 100 × CI in hibonites from CM2 chondrites; Ireland and Fegley 2000). This observation has been interpreted in terms of enrichment due to concentrating the REE in the earliest condensate, where 100 × CI would represent the first 1% of condensate from a cooling gas of solar composition (e.g., Boynton et al. 1980). In comparison, REE levels in most chondrules appear to be just above chondritic levels, suggesting lower temperatures for more complete condensation and sampling of the nebula in the chondrule-forming region.
Calcium-aluminum-rich inclusions show significant variability in REE abundances. This feature was used to classify them into five groups (Mason and Taylor 1982) (Fig. 4). Group I and V show overall flat abundance patterns, with the latter showing a small europium (Eu) depletion. Group III is depleted in both Eu and ytterbium (Yb), while Group VI is enriched in Eu and Yb. These inclusions show variability based around the most volatile of the REE: cerium (Ce), Eu, and Yb (Ireland and Fegley 2000). Group IV inclusions are olivine-rich, and are probably more closely related to chondrules than CAIs.
Of particular note are the Group II patterns. These show depletions in the most volatile REE (Eu and Yb), as well as the most refractory (gadolinium [Gd] through erbium [Er] and lutetium [Lu]). Group II patterns have a characteristic positive anomaly in thulium (Tm), which appears to have the volatility of the light REE. These patterns are interpreted as condensates formed after their separation from their ultrarefractory-enriched residue counterparts (e.g., Boynton et al. 1980). Both ultrarefractory-enriched, and ultrarefractory-depleted (Group II) CAIs are found, even in the same meteorite (e.g., Ireland 1990). These patterns appear quite complementary, suggesting a relatively close association of both residues and condensates during the high-temperature processing.
The volatility fractionation of REE is an important constraint on the thermal environment of the early solar system. The 50% condensation temperatures of REE range from 1360 K for Eu to 1480 K for Ce and Yb, through to 1580–1600 K for the other light REE and Tm, and up to 1660 K for the other heavy REE (at a pressure of 10−4 bar; Lodders 2003). The Group II pattern therefore suggests that it is produced by fractional condensation in the narrow temperature interval between 1660 and 1600 K. In this process, the first ultrarefractory-enriched material is separated from the gas, followed by continued condensation of lower temperature components.
It is noteworthy that major-element compositions of CAIs can be nearly identical, notwithstanding the notable differences in trace element abundances. This might indicate that major and trace elements are decoupled in some way, because the major elements cover a similar condensation temperature range as the trace elements and so wide variability could be expected in their abundances as well. Trace elements could have been fractionated in an early process, with subsequent mixing and melting with a more chondritic component that delivers the major elements.
Chronology of the Early Solar System
In the astronomical context, zero-age or t0, is generally associated with the ignition of the star, when equilibrium hydrogen fusion commences. In star formation, however, t0 usually refers to the onset of cloud collapse. Once ignited, a star will assume a position, largely defined by stellar mass, on the main sequence strip of the Hertzsprung-Russell (HR) luminosity-color (temperature) diagram. The time scale for stars to reach the main sequence is inversely related to their mass: a 0.5 solar-mass star requires about one billion years, whereas a solar-mass star needs 100 Myr (D’Antona and Mazzitelli 1994). In contrast, the preceding main accretion phase lasts up to a few million years (Calvet et al. 2000).
In meteoritics on the other hand, t0 is the age of the oldest objects, meaning that of CAIs. Meteoriticists can measure absolute and relative ages to extremely high precision (Amelin et al. 2002). Precise measurements of the ages of CAIs and chondrules have been obtained via the 235U-207Pb and 238U-206Pb coupled decay scheme. The U-Pb scheme gives an absolute framework to the age of the solar system because both U isotopes still exist and decay to their respective Pb isotopes. For uranium decay, both 235U-207Pb and 238U-206Pb can be used as chronometers. The 207Pb/206Pb ratio is also a chronometer if an initial 235U/238U is assumed or measured. Recently, it has been pointed out that the long-held assumption in geochronology of constant 235U/238U is incorrect for CAIs, and that this ratio must be measured independently (Brennecka et al. 2010). Amelin et al. (2010) report an age of 4567.2 ± 0.7 Myr for CAIs, where U isotopic fractionation has been determined independently rather than assumed as in previous measurements (e.g., Amelin et al. 2002). Similarly, Connelly et al. (2012) report chondrule ages with measured 235U/238U with ages that range from being contemporaneous with CAI formation through to around 3 Ma younger. It remains unclear whether chondrules represent a continuation of the process responsible for CAI formation, or whether chondrule heating occurs in a different spatial context to CAIs, or in a different spatial and temporal framework.
Other meteoritic evidence suggests that the CAI formation period may have been very rapid. Short-lived radionuclide abundances can be used as a chronometer, with respect to an inferred initial value of the parent/daughter abundance. For instance, 26Al (t½ = 0.7 Ma) decays to 26Mg, and the inferred abundance of 26Al (i.e., excess 26Mg) is commonly used as a chronometer. While 26Al has long since decayed, the abundance of 26Al relative to an initial abundance provides a relative temporal context. Thus, the ratio of 26Al/27Al in CAIs at 5 × 10−5 is taken as t0 for that system, and a lower inferred abundance of 26Al is taken as the signature of radioactive decay, thus allowing relative ages to be calculated. In addition, the dispersion of 26Al/27Al has been used as an indicator of the duration of the CAI formation event. Based on dispersion in 26Al/27Al of the main CAI population, the CAI formation epoch may have been as short as 4000 yr (Thrane et al. 2006; Larsen et al. 2011). There are two caveats to this interpretation. First, this is strictly a closure interval for that particular CAI, not a formation interval. The CAIs could have been above liquidus temperature for a longer period of time than is recorded by the uncertainty of the isochron. Second, CAIs have bimodal inferred initial 26Al/27Al ratios, with one peak at 5 × 10−5 and the other at less than 10−6 (MacPherson et al. 1995; Ireland and Fegley 2000; Krot et al. 2008). It could well be that different CAIs formed from regions with different 26Al/27Al, or formed at different times. As such, 26Al/27Al initial ratios should be used with caution in terms of interpreting time.
Winds and jets provide a dynamic scenario where grains could be heated and transported in the solar system. Liffman and Brown (1995) described a process of entrainment of particles in a protostellar wind, and estimated the size range of particles that could be ejected by an outflow powered by magnetic pressure (Liffman 2007) operating from the innermost regions of the disk (<0.1 AU). They obtained a droplet radius of less than 1 cm. Using a parameterized outflow solution, the authors found that 0.1 cm sized droplets can become decoupled from the outflow, and be ultimately incorporated into the accretion disk.
In a related scenario, also operating from the innermost disk regions, the X-wind model (Shu et al. 1996, 2001) proposed that a wind is accelerated centrifugally via the interaction between the accretion disk and the magnetosphere of the protostar. Solid material flowing toward the Sun from the protostellar disk is lifted by the X-wind and melted by solar radiation to form chondrules. The wind then transports the processed material outwards to the cooler environment of the outer solar nebula. While the X-wind model has attracted much attention in the meteoritics community, it does have some unresolved issues (e.g., Hezel and Palme 2008; Desch et al. 2010). Processing material close to the co-rotation radius can be problematic because in young protostellar systems the co-rotation axis lies within the dust sublimation radius (Flaherty et al. 2012), meaning that there is no dust at the launching point. This is in accord with observations suggesting that dust is indeed absent so close to the star (Muzerolle et al. 2003).
An important constraint from the meteoritics perspective is that if matrix and chondrules are complementary components of chondrites, then this suggests a close association of fine-grained matrix and large high-temperature inclusions throughout the chondrule-forming process (Hezel and Palme 2008). In the protostellar jet scenarios discussed above, chondrule and matrix thermal processing would need to take place in the vicinity of the protostar, followed by ejection back out to the accretion disk. Protostellar jets are an appropriate mechanism for transportation of processed material back out into the accretion disk. However, a model involving formation of chondrite constituents in this location would require mixing/unmixing of low- and high-temperature components of the matrix, accompanied by mixing of the high-temperature chondrules in the appropriate proportions. In such an energetic environment, maintaining these associations is difficult to envisage.
These issues encouraged us to look at disk-winds as a potentially suitable mechanism for the thermal processing of chondrule material. In a quiescent protostellar disk, particles are expected to settle to the midplane (e.g., Dullemond and Dominik 2005) (see also Fig. 5). However, protostellar winds carry enough momentum to lift dust grains via collisions with the outflowing gas (Safier 1993). In Salmeron and Ireland (2012; henceforth SI12) we showed that a disk-wind operating at a radial distance of 1 AU can thermally process dust particles embedded in the outflowing gas (see Königl et al.  and Salmeron et al.  for details of the wind models). We found, specifically, that very small particles (with an initial radius of about 0.01 cm) are rapidly accelerated away from the disk by the outflow. More massive, centimeter-sized particles would fall to the midplane, as the gravitational force consistently dominates over the drag exerted by the outflowing gas. There is, however, an intermediate size regime (particles with initial radius approximately 0.05–0.1 cm) for which the grain is initially lifted by the wind but, as it grows, the downward pull of gravity becomes dominant and causes the particle to turn back toward the disk (see Figs. 5 and 6). In one model, an initially 0.1 cm sized particle rose up to about 0.05 AU above the disk, where it had grown to about 0.15 cm, before turning downward.
The temperature of the grain changes in response to interactions with the gas and with the external radiation field. In SI12 radiative equilibrium is assumed, so that the grain’s radiative losses are balanced by radiation absorbed from adjacent particles (Hood and Horányi 1991). As the grain moves upwards, the gas-dust drift velocity remains small for most of its trajectory, as both gas and particles accelerate. However, as the grain grows and becomes heavier, it slows down and turns downwards, causing the dust-gas relative velocity and associated heating to increase dramatically. In SI12, a particle with an initial size of 0.05 cm reaches approximately 0.2 AU above the midplane, where it has grown to 0.08 cm, and experiences a maximum temperature of about 2000 K, both typical values of chondrule material. The grain then cools and falls back to the disk, where it may reassemble with the relatively cool material of the disk interior.
Disk-Wind Processing and Basic Chondrule Properties
SI12 showed that chondrules with an initial temperature of about 400 K (consistent with the typical temperatures at AU distances inferred for protostellar disks; e.g., D’Alessio et al. 2006) can be melted in a disk-wind at a distance of about 1 AU from the proto-Sun. Moreover, the initial temperature of the dust has little effect on the ultimate temperature experienced by the chondrule. Peak temperatures of 1800 K were reached in the SI12 models for chondrules of size around 0.1 mm. Cooling occurs when the particle is removed from the fast wind as it turns downwards. SI12 obtained maximum cooling rates of order 1 K h−1 assuming continued radiative equilibrium. However, they also noted that cooling rates could be much faster if radiative cooling is incorporated in the models, as this effect would be important once the particle rises to the optically thin atmosphere.
SI12 proposed that formation in a wind-driving disk can naturally explain a number of basic chondrule properties (e.g., Ciesla 2005; Hewins et al. 2005; Scott and Krot 2005b; Hezel and Palme 2008; Krot et al. 2009): They are typically within a size range of 0.1–1 mm, which is a natural consequence of the processed particles being small enough to be lifted up by the outflow, and large enough not to be ejected from the disk altogether. The high proportion of chondrules in chondrites reflects the efficiency of processing material in the region of the disk where the wind operates. In such an efficient process, compound chondrules would be expected from collisions of similarly sized precursors at relatively low velocity. Moreover, the more massive compound chondrules would be expected to settle back to the midplane faster than individual chondrules, and therefore possibly cool faster as well. During settling, it would also be expected that chondrules would accrete dusty rims from the cold environment of the disk.
An attractive aspect of a disk-wind scenario is the possibility of preserving an association of the cold material in the disk with the chondrules that are forming in the jet. This leads to the possibility of explaining, to some degree, the complementarity observed between matrix and chondrules. The initial heating of material will liberate volatile elements and these will either be entrained in the wind and carried away, or some fraction of this material may interact with the disk and remain available for reassembly with the chondrules.
A disk-wind origin is also quite compatible with oxygen isotopic properties of chondrules in a broad sense. Oxygen isotope compositions are generally expressed as δ17O and δ18O, which are the fractional differences in oxygen isotope ratios relative to terrestrial expressed in parts per thousand. Chondrules typically show oxygen isotopic compositions that range from small deficits in 16O relative to terrestrial (approximately δ17O ≈ δ18O approximately +10‰), to moderate enrichments in 16O (δ17O ≈ δ18O down to −20‰) (e.g., Yurimoto et al. 2008). One chondrule is noteworthy in having the largest enrichment in 16O yet observed (δ17O ≈ δ18O = −80‰; Kobayashi et al. 2003), and thus appears even more 16O-rich than the Sun (at δ17O ≈ δ18O approximately −60‰; McKeegan et al. 2011) and refractory inclusions with approximately the same extreme composition.
The isotopic composition of solar system materials can be interpreted in terms of the photodissociation and self-shielding model of Yurimoto and Kuramoto (2004), with inheritance of the oxygen isotopic signature of the parental molecular cloud. In this model, CO is dissociated by UV photons, producing atomic C and O. UV self-shielding occurs because the column density of C16O is much greater than that of C17O or C18O. Thus, in cloud interiors, dissociated 17O and 18O react with abundant hydrogen to form 17O-18O-rich H2O, and residual C16O becomes relatively more abundant (with respect to total CO). The initial refractory dust in the solar system, with close-to-solar composition, then reacts with nebular water to form a planetary reservoir (i.e., close to terrestrial) that is enriched in 17O and 18O relative to solar.
As a result of the process outlined above, the disk-wind associated with the solar nebula could assume a composition that is a mixture between the CO, planetary dust, and nebula-water oxygen isotope compositions (Ireland 2012). If all components are present in solar proportions then the gas will have a solar composition, i.e., 16O-rich relative to terrestrial. If the dust/gas ratio increases during chondrule formation as a result of dust settling and growth processes, as is expected (D’Alessio et al. 2001, 2006), then the composition will be closer to the planetary reservoir. Further variability would be produced by addition of 16O-rich CO or 17O-18O-rich nebula water. Thus, the chondrule composition will reflect a specific component mix, which could be region-specific to the solar nebula, and therefore variations in oxygen isotopic composition could be a natural consequence of the temporal and spatial evolution of the disk.
The distribution of oxygen isotopic compositions of chondrules suggests that they mainly formed in an 17O-, 18O-rich nebular gas (Yurimoto et al. 2008). The oxygen isotope composition of the gas in the chondrule-forming region therefore appears to be close to the planetary composition, either through the evaporation of pre-existing silicates that had already reacted to this composition, or through an appropriate mixture of the primordial components. Chondrules heated to melting point in a disk-wind would quickly assume the oxygen isotopic composition of the ambient gas.
Such a scenario also has consequences for the chemical oxidation state of chondrules. Chondrules are divided into two groups depending on the fayalite (Fe2SiO4) content in olivine ([Mg, Fe]2SiO4). Type I chondrules have a fayalite component of less than 10%, whereas Type II chondrules have fayalite contents higher than this value. Grossman et al. (2008) document the problem of high Fe2+ contents in olivine in a gas of solar composition, which is highly reducing. However, the requirement for a highly oxidizing nebula gas can also be placed in the context of a disk-wind whose composition depends on the ambient gas as well as on the evaporating species of the disk. In a reducing wind with abundant H and CO, Fe will be present as metal and will not be compatible with the olivine matrix. If the wind is more oxidizing, due to admixture of nebula water and evaporated dust, Fe will move to the Fe(II) state where it is compatible with the olivine matrix. Gas/dust ratio and size of grains may also be relevant in this phenomenon. A natural consequence is that oxygen isotope compositions of Type II chondrules with oxidized iron could be heavier than Type I chondrules if the oxidation is a result of the presence of the 17O, 18O-rich nebula water for example. However, this situation is unlikely to be as clear cut as this and indeed Rudraswami et al. (2011) show that while Allende Type I chondrules can extend to quite 16O-rich compositions, they show a wide range in oxygen isotope compositions (Δ17O from −18 to +1‰), whereas Type II chondrules have a more restricted range (Δ17O from −6 to −2‰; Rudraswami et al. 2011).
One of the great enigmas of chondrule formation is that they have been taken to temperatures sufficient for melting olivine, yet they retain relatively volatile elements such as sodium (Na). Alexander et al. (2008) have shown that at least some chondrules formed essentially as a closed system with respect to Na. This suggests that the vapor pressure of Na must be very high, and can be equated to a gas cloud incorporating 0.01–0.1% dust.
The high dust-gas ratio might also be a constraint on the degree of isotopic mass fractionation that the chondrules experience. The higher above the liquidus the temperature rises, the shorter the period of time they can stay at that temperature before magnesium (Mg) and silicon (Si) begin to evaporate, producing isotopic mass fractionation if surrounded by vacuum. Therefore, the limited isotopic mass fractionation effects in Mg and Si suggest that chondrules did not evaporate extensively to vacuum, but rather that evaporation took place in a nebula containing a substantial dust component, and/or that the chondrules were not heated significantly above the liquidus temperature, and/or the heating lasted only a short time. The lack of potassium (K) isotopic fractionation also indicates little loss to vacuum. K is relatively volatile (50% condensation temperature of 1000 K) and so is expected to have evaporated at the melting temperature of olivine.
In a disk-wind, with a relatively large dust-to-gas mass ratio (the canonical figure is 1%), satisfying the above conditions is not particularly problematic. Small particles would be evaporated relatively quickly, enriching the nebula gas in moderately volatile rock-forming elements. As such, the lack of isotope fractionation could be explained by the high effective partial pressure of rock-forming elements.
Note that in Fig. 6 the peak temperature attained by any given grain is size dependent. This trend may be an artifact of our modeling procedure, as we have treated the grains during the modeling as being entirely refractory. In reality, the grain will begin to melt when the solidus temperature is reached and the temperature rise will be stalled as melting proceeds. The molten droplet will then behave quite differently in the jet, potentially spalling and losing mass, or gaining mass by spattering of small droplets. Ultimately it appears that all chondrules have become molten with limited evaporation.
Disk-Winds and CAIs
A key question regarding the origin of our solar system is whether CAIs formed in the same type of process as chondrules, or whether a different process is required. In the context of this discussion, CAIs could also be a product of disk-winds operating in the solar nebula, possibly in a different location, e.g., closer to the Sun, or at a slightly earlier time.
Conceptually at least, this appears to be a reasonable possibility. Liquidus temperatures of CAIs are generally lower than chondrules, so the peak temperature is probably not an issue. The main distinction between CAIs and chondrules is the presence of Mg and Si isotopic fractionation in CAIs, especially in the FUN (fractionated and unidentified nuclear effects) CAIs, and in some related hibonite grains, where extensive isotopic mass fractionation—and by inference evaporation—is apparent in many elements that are regarded as refractory.
We have already discussed that disk-winds can heat a millimeter-sized chondrule to around 2000 °K. CAIs require similar temperatures for condensation from a gas of solar composition, or for melting, but could require higher temperatures for evaporation and fractionation of isotopes. However, it is not simply a high temperature that is required. The largest CAIs (approximately 20 mm in size) are an order of magnitude larger than the particles we have modeled thus far. To float these objects, much greater gravitational forces must be overcome. A possible scenario would be a more energetic wind located in closer proximity to the Sun, in agreement with the observed onion-like structure of protostellar jets where faster material is being ejected from smaller radii (e.g., Cabrit 2007a). Alternatively, particularly energetic winds may be associated with an outburst state of the young star, such as the FU-Orionis and/or EXor phenomena. It is critically important to explore the parameter space where disk-winds operate, and the dependence of the disk and wind parameters with the size and peak temperatures of processed precursors, to ascertain whether winds operating closer to the Sun, and/or during an active state of the protostar, would be strong enough for the thermal processing of CAI material.
Ultimately, it appears that CAIs cannot be the result of a single, simple condensation event in the solar nebula (Ireland and Fegley 2000). The volatility fractionation of REE occurs at extremely high temperature, and well above the condensation temperatures of Mg and Si. While CAIs with the Group II condensate pattern could be produced by continued condensation, the ultrarefractory-enriched patterns cannot have a simple single-stage condensation origin. They must have been isolated from the thermal processing very early on, prior to full condensation of REE and other refractory trace elements. However, the major-element chemistry of CAIs is quite similar no matter the trace element abundance pattern. A scenario for CAI formation must involve evaporation and condensation of refractory elements at high temperature to produce dust/gas with a range of volatility fractionations. This dust-gas is then mixed with Ca-Mg-Al-Si-O-rich dust and the proto-CAIs are raised in temperature to the melting point to form compact objects. Higher temperatures would lead to evaporation and condensation of “fluffy” CAIs.
The disk-wind provides a setting where evaporation and condensation are simultaneously active processes. Small particles are likely to be carried away or completely evaporate and thus will have a limited contribution to forming a Group II pattern. Larger particles will experience extensive evaporation; however, an ultrarefractory residue could survive. Therefore, the gas above this region could have a signature that is depleted in the ultrarefractory trace elements (i.e., Group II). This gas will ultimately condense, forming refractory dust which is accreted onto other dusty objects with a more chondritic chemistry. Heating of these objects as they move through the jet could then lead to melting and evaporation, forming compact (melted) and fluffy (condensate) CAIs.
The new chronology appearing for CAIs and chondrules is intriguing. Contemporaneous production of CAIs and chondrules possibly indicates their formation in different regimes of the solar nebula. In terms of thermal processing, it is tempting to advocate that CAIs were only produced in the highest temperature locations close to the Sun and these conditions only lasted for a short time. The continued production of chondrules could then reflect the motion of the appropriate thermal regime for chondrule production closer to the star as it evolves. This picture would be consistent with the presence of 10Be-10B anomalies in CAIs with 10Be being produced by spallation from an early active Sun (McKeegan et al. 2000). If chondrules formed outside this inner zone, then their exposure to solar radiation could have been more limited.
However, as has been observed many times before, there remain distinct differences between CAIs and chondrules. The preservation of extremely large isotopic anomalies in 48Ca and 50Ti in CM2 hibonites, and the smaller but ubiquitous anomalies in CV3 CAIs indicate that the formation region of the CAIs was distinct in preserving these components. Why chondrules do not see this heterogeneity is perplexing. Even more so if we consider that Ti isotopic anomalies appear to be still present while chondrites were accreting (Trinquier et al. 2009). Now that we see chondrules with similar ages to CAIs, it is important to see if there are any differences in early versus late chondrules, and whether the early formed chondrules preserve any isotopic anomalies comparable to CAIs.
As noted above, oxygen isotopes differ between CAIs and chondrules, but this may simply reflect a different gas composition (i.e., ratio of “solar” to “planetary” components; Ireland 2012). But this gas composition could also have a radial structure in the solar nebula with only CAIs seeing the bulk solar gas rather than a gas composition dominated by evaporating silicates and oxides. Again, a correlation between chondrule ages and compositions could elucidate the structure of the early solar nebula.
SI12 concentrated on the most general chondrule properties, as that work was intended to show a “proof of principle,” that disk-winds are a suitable mechanism for chondrule formation. On the other hand, there are important properties of chondrules that were not contemplated by that study. These are the many textural, petrological, and isotopic characteristics of this material that contribute to their diversity and complexity. A number of these properties have been discussed above. Below we point to some important avenues for further research in the context of the disk-wind model for chondrule processing.
Radial Extent of the Wind-Driving Region in Protostellar Disks
The extent of the wind-driving region in protostellar disks is not well constrained; however, current observational evidence seems to indicate that it extends to a few AU from the central protostar (e.g., Coffey et al. 2004, 2007; Agra-Amboage et al. 2011). Studying the viability and properties of protostellar winds as a function of radius from the protostar, and the processing of dust at such distances, is a distinctly fertile area of research that may, in turn, help establish the extent of the chondrule-formation region. It is also interesting to study the efficiency of the chondrule-formation process at different radial distances in the disk, and links between the radius of formation and the resulting chondrule properties. Interestingly, should this mechanism be viable at distances of order approximately 5 AU, then it could be possible to explain the presence of thermally processed material in comets without having to invoke the action of a long-range transport mechanism to bring them to such distances.
Chondrule/CAI Formation during Protostellar Outbursts
Young stars are often seen to experience dramatic outbursts, during which the optical brightness of the source increases by a factor of a hundred or more (e.g., see the review by Hartmann and Kenyon 1996; and Aspin 2011). These outbursts are thought to be an integral part of the star-formation process. They are thought to be associated with a significant increase in the accretion rate onto the central object, and therefore, represent stages during which a substantial fraction of the stellar mass may be gained. Two types of outburst activity have been found. FU-Orionis-type outbursts are the most energetic, with the source brightness increasing by over five magnitudes and lasting for approximately 100 yr. The accretion rate of these sources is approximately 103–104 that of the source in quiescence (Hartmann and Kenyon 1996). The EXor-type outbursts result in an increase in brightness of 1–4 magnitudes and 100-fold jump in the accretion rate (e.g., Aspin 2011) during a timescale of weeks to months (up to about 2 yr in particularly energetic cases). It is interesting to explore chondrule formation in these environments, and particularly, the formation of CAIs in the most powerful winds associated with these sources.
Thermal Processing in Disk-Winds with a More Detailed Treatment of Radiation
In SI12 we assumed that a population of small particles (assumed to be significantly smaller than the chondrule precursors) is suspended with the fluid and thermally interact with the precursors. We further assumed that the energy absorbed by the chondrule as a result of the emission of neighboring grains balances its own radiative energy losses (e.g., radiative equilibrium holds). This is a commonly adopted approximation in models of dust-rich environments; however, it is expected to break down in low-density, optically thin regions. Adopting this approximation is particularly problematic for the case when all the grains have the same size as the precursors. In this case, a dust-to-gas mass ratio enhancement of up to a factor of ten over the canonical value (0.01) is required for it to hold at elevations beyond approximately 1 scale-height from the disk midplane (SI12). On the other hand, when the surrounding dust is assumed to be 0.01–0.001 cm in size, radiative equilibrium holds for the entire trajectory of a precursor with an initial size of 0.05 cm when the dust-to-gas mass ratio is 0.03–0.01, respectively (SI12). It is worth noting here that although at the midplane the density scale-height of wind-driving disks is significantly smaller than the tidal scale-height (as a result of magnetic compression; Wardle and Königl 1993), at higher elevations this trend reverses, so the drop in density with height is substantially less pronounced than in quiescent disks. Studying the processing of chondrules in disk-winds incorporating a detailed treatment of the dust radiation properties is critical to refine the estimates of the peak temperatures attained by the precursors as a function of size, as well as the cooling rates they experience.
Protostellar Disk Background Temperature Structure
SI12 assumed that the background temperature of the disk was independent of height, and consistent with the inferred midplane value of an actively accreting disk at a radial location of 1 AU. Under these conditions it was found that the peak temperature of the chondrule precursors was not strongly dependent on the assumed gas background temperature (which is also the initial temperature of the dust). It is, however, known that in the disk atmosphere the gas may become significantly hotter than the dust, as a result of gas heating by stellar irradiation and internal dissipative processes, together with weak gas-dust thermal coupling (e.g., D’Alessio et al. 1998; Glassgold et al. 2004). Glassgold et al. (2004), in particular, found that the gas temperature high above the midplane could potentially reach thousands of degrees, whereas the dust temperature remains below approximately 500 K. It is important to explore the effect of an elevated background temperature, as well as its vertical profile, on the processing of chondrule material. Given that the recovery temperature (which influences the amount of dust-gas heat exchange) increases with the background gas temperature (see Hood and Horányi 1991), it could, in turn, affect the heating experienced by the chondrule precursors.
Chemistry and Implications for Grain Properties
SI12 followed the evolution of only a few grains, with particular reference to their temperature history and size. As indicated above, a disk-wind may provide a setting where the diverse chemistry of CAIs and chondrules could be explained. For instance, sodium concentrations in chondrules are known to be orders of magnitude higher than expected (Alexander et al. 2008). This property, along with the lack of potassium isotope fractionation, suggests that processing of the precursors took place in a gas with an extremely high concentration of these elements. Such a scenario is difficult to produce in a closed system. However, a disk-wind could provide elevated concentrations of these elements due to volatilization at lower elevations, thereby stabilizing the alkali metals in chondrules.
The astrophysical setting of the early solar system is becoming better constrained through the combination of astronomical observations, theoretical/numerical modeling, and experimental studies of the meteoritic samples. However, taking account of the myriad of details obtained from meteoritic analyses in this overall picture is problematic, because of the extreme differences of scales between astronomical and meteoritical studies. Furthermore, the meteoritical evidence concerning the conditions during the birth of our own solar system 4.57 billion years ago is seemingly at odds with the astrophysical view, in that a large fraction of material throughout the early solar system seems to have experienced high-temperature processing. In contrast, astronomical evidence suggests a relatively cool accretion-disk environment, with high-temperature processing likely only in close proximity to the young stellar object.
Protostellar winds accelerated magnetically from the disk surfaces provide a plausible scenario for the thermal processing of material, which is consistent with the transient nature of chondrule and CAIs formation and the maintenance of an overall cool nebular environment. We have demonstrated that heating of chondrule-sized particles out to 1 AU is possible, but disk jets could operate closer to the Sun as well as farther out.
Acknowledgments— We are grateful to Kurt Liffman and Martin Bizzarro for their thoughtful reviews and to the associate editor Sasha Krot for comments and advice. This work was supported by the Australian Research Council through DP120101792.