Setting the Stage: The Early History of the Solar System

Abstract

This article reviews the early history of our solar system from an astrobiological perspective and presents evidence from meteorites and astronomical observations. The purpose is to trace the formation of key molecules that participated in the building blocks of life. The Sun and its planetary system started from a section of a molecular cloud that collapsed into a protoplanetary disk. In the center of the protoplanetary disk, the protosun heated the surrounding material. The dust and gas inherited from the cloud remained pristine farther away from the protostar, while new compounds were created in the gas and on the icy mantles of the dust. The dust accreted into pebbles, pebbles formed planetesimals, and planetesimals collided and accreted pebbles to create planets. Meanwhile, the protosun became the Sun when its core reached the pressure and temperature required to transform hydrogen into helium. During this process, the Sun emitted high-energy radiation and particles that impacted the chemistry in the disk and the early evolution of the terrestrial planets.

Keywords

Prebiotic chemistry—Protoplanetary disks—Solar system—Protosun

1. Introduction

During the first half of the last century, scientists recognized that meteorites and solar system formation were linked to the earliest composition of Earth’s atmosphere; this provided the first geological context for the origins of life (Bernal, 1949; Haldane, 1929; Oparin, 1938; Urey, 1952). At that time, we did not have a clear understanding of how our solar system was formed. But today we can trace our cosmic origin back to the Big Bang event about 13.8 billion years ago, which led to the creation of hydrogen atoms, and continued with the first stars that formed carbon, nitrogen, and oxygen. For the purpose of this review, we begin in Section 1 with the formation of our star and end with the building of the planetesimals that resulted in our planet. These planetesimals further provide clues to the pathways that led to the origins of life. Section 2 is dedicated to the formation of the solar system based on two lines of evidence: young planetary systems and bodies in our system, which include asteroids, comets, and meteorites; here we focused on the chemical properties of these objects to trace the path of the main carriers of the essential elements for life. Section 3 presents the characteristics of the early Sun, its evolution in time, and the impact of early solar activity on the environment that led to the origins of life.

2. The Early Solar System

The history of the solar system has been constructed from several lines of evidence. The geological record provides limited constraints on the earliest stages of our planet. The majority of our understanding of the early solar system comes from observations of young stellar systems, comets and asteroids in our solar system, the remnant pieces of asteroids that fall to Earth as meteorites, laboratory experiments, and theoretical simulations. In the last decade, we have added to this evidence the architecture of other planetary systems that serve as probes of our planet formation scenarios, and samples of asteroids and comets analyzed in situ or collected and brought to Earth to be studied. We will focus mainly on tracing the chemical compounds that may have contributed to the origins of life and the establishment of a habitable planetary environment.

2.1. Solar system formation: Evidence from young stellar objects and protoplanetary disks

Stars and planets are formed in molecular clouds, the densest and coldest regions of the interstellar medium (ISM), with densities of H atoms of n_H = 100–10⁵ cm⁻³ and temperatures of T ∼10 K (Heyer and Dame, 2015). Molecular clouds are formed of H and He and the accumulation of the remnants of stars at their latest stages. The formation of molecular clouds is the result of thermal, dynamic, and gravitational processes that preclude the cloud from heating or expanding at large scales (10–100 pc, Ballesteros-Paredes et al., 2020). These clouds are comprised of gas and dust, and to date, more than 300 molecules have been identified in the ISM (Müller, 2024). In molecular clouds, the most relevant molecules after hydrogen (H₂) are carbon monoxide (CO), water (H₂O), and organic molecules mainly composed of H, C, and O (Agúndez and Wakelam, 2013; Bergin and Tafalla, 2007). The dust is composed of silicate grains and/or hydrogenated amorphous carbon that can be covered by CO, H₂O, methane (CH₄), and ammonia (NH₃) ices, depending on the environment; other possible components are graphite and nanodiamonds (more details in Section 2.3.1).

Pre-stellar cores are the densest objects in molecular clouds and are typically associated with filamentary structures (Hacar et al., 2023), with sizes < 0.1 pc and densities n_H > 10⁵ cm⁻³ (Fig. 1). The cores can become gravitationally unstable, which can lead to the formation of one or more young stellar objects (YSOs) (Rosen et al., 2020), also referred to as protostars or premain sequence (PMS) stars (Tobin and Sheehan, 2024). This is the evolutionary stage before the onset of hydrogen-burning nuclear reactions, which produce energy in the core of the star. The star enters the main-sequence stage at the onset of hydrogen fusion. Young stellar objects are typically classified by their infrared (IR, 2–25 µm) spectral energy distributions (wavelength-dependent emission), which provide a comprehensive timeline for stellar formation starting from the core to the planetary system stage (Lada, 1991; Williams and Cieza, 2011). The classification of YSOs initially included three categories, Classes I, II, and III (Lada, 1987), from younger to older objects (Fig. 2). Later, a Class 0 was added to include strong submillimeter emission cores that had density profiles consistent with central internal heating (André et al., 1993).

FIG. 1.

Cartoon of the hierarchical structure of a molecular cloud, not to scale (not to scale, after Pokhrel et al., 2018). Diffuse ISM can have atomic H or H₂ with densities of <500 atoms of H/cm³ (Snow and McCall, 2006). Cores have sizes from 0.05 to 0.1 parsecs (pc) and envelopes from 0.0015 to 0.015 pc (Pokhrel et al., 2018).

FIG. 2.

General chronology of the early solar system. The upper images show examples of young stellar objects in key evolutionary classes, from most embedded to most evolved. The timeline illustrates important processes in the solar protoplanetary disk traced from asteroids and meteorites and their components (e.g., Halliday and Canup, 2022). The gradient in color bars shows uncertainties on the ages of objects or processes with lighter shades in a color bar correlating to greater uncertainty. Pebbles (mm-cm-sized particles) are accreted from dust (micron-sized particles) and are generated by impacts of larger objects. Images (upper panels, left to right): Barnard 68 (ESO, https://www.eso.org/public/images/eso9924a/), HH30 (Tazaki et al., 2025), TW Hydrae (Andrews et al., 2016), NO Lup (Fig. 2 in Lovell et al., 2021), and HD170773 (Sepulveda et al., 2019).

Class 0 objects largely emit in the far IR and millimeter wavelength ranges, with temperatures derived from their IR emission of T ∼15–30 K (André et al., 2000). They are considered molecular cloud cores or sections of a core collapsing into a “flat” structure known as a circumstellar or protoplanetary disk. Disks begin to form during the Class 0 stage on timescales of ∼10⁴ years (Williams and Cieza, 2011). At this stage, the protostar is surrounded by a disk and an infalling envelope (see Fig. 1) with a typical size of ∼1000 astronomical units (au, with 1 au being the average Earth–Sun distance) (Rosen et al., 2020). The dust in the disk and the infalling envelope is responsible for the IR and millimeter emission, which fades as the YSO evolves to a star, while the envelope is dispersed by the radiation of the central YSO as the disk evolves. One of the main characteristics of Class 0/I objects is the molecular outflows that can be highly collimated (jets) or conical winds with velocities on the order of 10–100 km/s with lengths of < 1 pc (Bachiller, 1996; Bally, 2016).

Young stellar objects transition from Class 0 to I when the protostar heats the envelope to a temperature of T∼70 K (Rosen et al., 2020). The Class 0 stage lasts ∼ 0.1 Myr, and Class I ∼ 0.5 Myr (Rosen et al., 2020). At these early stages, the disks are embedded in opaque envelopes; therefore, their masses and radii are poorly constrained even with dedicated observations with instruments such as the Atacama Large Millimeter/submillimeter Array (ALMA), which can achieve high spatial resolution (tens of milliarc seconds) (Andrews, 2020; Tung et al., 2024). Once the envelope disappears, the YSO is classified as Class II or a classical T-Tauri star, although it has been noted that these two classifications may not be an exact match (Williams and Cieza, 2011). The Class II YSO is the PMS star with a protoplanetary disk. Stars with masses between ∼0.1 and 1 M_⊙ have disk masses that range from ∼0.01 to 0.1 M_⊙ and radii that extend to hundreds of astronomical units (Williams and Cieza, 2011), with a typical dust-to-gas mass ratio of 1:100 (Zhao et al., 2020).

The evolution of protoplanetary disks is largely dictated by the rate of material that accretes onto the disk and protostar, dust growth to larger bodies, dynamics between those bodies and the star, and photoevaporation that eventually leads to the total loss of gas. Class II objects have lifetimes of 1–3 Myr, after which they transition to Class III objects, which is when their disk gas is mostly lost, and dust becomes transparent in the IR (Williams and Cieza, 2011). Class III disks are harder to identify in the IR due to their low IR excess. Therefore, they often require millimeter and submillimeter observations, and they can be further identified by the X-ray characteristics of their PMS star (e.g., Bontemps et al., 2001; Williams et al., 2019). During the Class III stage, the dust mass drops to < 0.1 Earth masses (M_⊕) and becomes a debris disk (Manara et al., 2023). The gradual transition between the Class III and debris disk is observed in young clusters with ages < 5 Myr, where the dust of many Class III objects has overlapping properties with debris disks, which in turn implies very short timescales for planet formation (Manara et al., 2023 and references therein). Class I to Class III objects are key to our understanding of the processes that lead to planet formation and will be reviewed in the next section.

2.1.1. Protoplanetary disks and planet formation

Here, we describe the general process of terrestrial planet formation derived from observations and theoretical models. Protoplanetary disks are initially made from dust and gas reservoirs that are either inherited from the molecular cloud or reprocessed in the envelope of the YSO, while the densities and temperatures change during the infall of the envelope (see Section 2.3.1 for details on the composition of disks). Planet formation can generally be described as the growth of dust (microns) to larger objects (hundreds to thousands of kilometers), with processes for each scale that in turn depend on the disk pressure, temperature and density profiles, and chemical composition. One of the most relevant features of the disk is the water snowline, the distance from the protostar where the temperature is low enough (∼100–150 K) to allow the condensation (the process of converting gas to solid) of water onto the grains. Inside the snowline, dust is dominated by silicates, and terrestrial planets should grow within this radius. The position of the water snowline moves depending on the temperature profile of the disk; for example, the model by Chambers (2009) predicts that the water snowline can move from beyond ∼7.5 au (Class 0, age 40 kyr) to ∼1.7 au (Class II, 2 Myr) for a solar-mass star. However, during the Class 0 and I stages, the location of the water snowline can move closer or far from the YSO during transient events that heat the disk, such as sudden increases in the accretion rate or flares (Krijt et al., 2023).

The initial steps of planet formation involve dust grains that grow from diameters of a few micrometers to centimeters, driven by electric forces, acquired by mechanisms such as tribocharging (surface electrification by contact), that allow the grains to stick together, a process facilitated by the low relative velocities of the grains (Birnstiel, 2024; Onyeagusi et al., 2025). Dust grains smaller than ∼1 mm are anchored to the gas and follow a Brownian motion, while larger grains are dominated by turbulence (Birnstiel, 2024). The interaction between the gas and particles becomes hydrodynamic for sizes >0.1 cm due to the differential velocities between the dust that orbits the star following a Keplerian velocity (proportional to r^−1/2) and the gas that orbits at a sub-Keplerian velocity due to its own pressure. The dust is decelerated by the gas; i.e., it loses momentum that causes it to drift toward the star (Birnstiel, 2024; Weidenschilling, 1977). Then, as the grains drift toward the star, the grains grow and settle at the mid-plane of the disk. Grains stop growing when the timescale of radial drift is larger than that of collisions, with the maximum grain size depending on the distance to the star (Birnstiel, 2024). Collisions can result in growth, bouncing, or the generation of smaller particles either by fragmentation or erosion (Güttler et al., 2010; Zsom et al., 2010). Recent models suggest that the maximum sizes are larger (up to centimeters) at smaller disk radii (∼1 au) (Andrews and Birnstiel, 2018; Birnstiel, 2024); consequently, centimeter-sized grains drift toward the star without growing and therefore prevent planet formation. Thus, radial drift is a barrier for planet formation for smooth-density-profile disks. However, if pressure bumps (disk zones with a local pressure maximum) form, they can effectively trap grains which prevents further radial drift of the dust. Pressure bumps may be created, for example, by vortices or by the presence of a growing giant planet. Although the mechanisms are not yet fully understood, there is observational evidence that shows disks with structures which would result in local pressure maxima, such as rings and spiral arms (e.g., Andrews, 2020, Fig. 11). While dust drifts and settles, the relative velocities of the grains increase, and subsequent grain–grain collisions lead to fragmentation and bouncing, which then become barriers to further grain growth (Güttler et al., 2010; Zsom et al., 2010). Collectively, these phenomena are termed the “meter-size barrier” (Birnstiel, 2024; Krijt et al., 2023).

There are two current hypotheses for the formation of kilometer-sized objects from millimeter to centimeter grains or “pebbles.” The first hypothesis centers on the process of coagulation, where larger objects are formed by collisions of smaller ones. The main problem with coagulation is the fragmentation and bouncing that prevent object growth (Drazkowska et al., 2023; Johansen et al., 2014; Simon et al., 2024). Recent models show that highly porous icy particles or silicate nano-grains can grow to sizes of 100 m in timescales faster than radial drift, although other theoretical models suggest that this mechanism is unlikely (Drazkowska et al., 2023; Johansen et al., 2014; and references therein). The second, and leading, hypothesis relies on streaming instabilities that result from a runaway process where a dust overdensity is fed by radial drift and Coriolis forces (Simon et al., 2024 and references therein; Youdin and Goodman, 2005; Youdin and Johansen, 2007). The clumps formed by this mechanism must surpass the Roche density to produce gravitationally bound dust clouds that will become ∼100 km objects, or planetesimals (Watanabe et al., 2023).

Once planetesimals are created, the growth to planets proceeds by collisions, which are aided by the self-gravity of the planetesimals (Kokubo and Ida, 2000,1998). This process halts when the planetesimals available to enter the Hill sphere of a planetary embryo are consumed. The mass reached at this point is called the isolation mass; for a solar-mass star, the isolation mass is ∼0.1 M_⊕ at around 1 au. Planetary embryos outside the water snowline can grow up to 10 M_⊕, which allows the embryo to accrete disk gas, which leads to the formation of gas giants (Drazkowska et al., 2023). Planetary growth may be aided by pebbles dragged by gas toward planetary embryos (Johansen and Lambrechts, 2017; Ormel, 2017) and embryo-embryo collisions (Agnor et al., 1999; Elser et al., 2011; O’Brien et al., 2006). Pebble accretion is currently thought to be most relevant for exoplanets with masses or radii larger than Earth formed at or beyond the water snowline (Bitsch et al., 2018; Youdin and Zhu, 2025). For the solar system, several lines of evidence point toward collisions between planetary embryos as the main process that determined the final size and mass of Earth, and probably Mercury and Mars (Bizzarro et al., 2025; Morbidelli et al., 2025), although this matter is still debated (Johansen et al., 2014). Planetesimals, planetary embryos, and newly formed planets undergo dynamical interactions with one another and with the disk gas. Consequently, planetary embryos and planets migrate, which perturbs the orbits of planetesimals and sends material from outside the water snowline to closer radii from the star (Raymond and Izidoro, 2017: Nesvorný, 2018) and, in some cases, expels planets from the system (Carrera et al., 2019; Ford and Rasio, 2008; Jurić and Tremaine, 2008).

Although all planetary systems undergo the same processes, their final architectures depend on several factors (Raymond et al., 2020, and references therein):

1.
The mass of the disk, where a more massive disk accelerates the formation of planetesimals, planetary embryos, and gas giants, which can lead to larger giant planets;
2.
The time and place of planet formation, which determines the sizes of the resulting planets and the location and composition of the final debris;
3.
The formation of giant planets, which prevents the accretion of pebbles onto bodies within the orbit of the giant planet embryo and blocks the inward migration of smaller planetary embryos;
4.
Dynamic instabilities induced by giant planets that promote the growth of terrestrial planets by increasing the likelihood of planetesimal impacts.

The solar system has a unique architecture compared to other planetary systems thus far observed. For example, the most common exoplanets observed thus far have radii between 1 and 4 times Earth’s radius (R_⊕) (Winn and Petigura, 2024), sizes that do not exist in the solar system (Raymond et al., 2020). Also, numerical models predict a larger mass for Mars, with one possible proposed solution being a depletion mechanism in the zone where Mars was formed. Another recent scenario points to all terrestrial planets forming from a ring of planetesimals with a mass of ∼2 M_⊕ (Raymond et al., 2020, and references therein). Meteorites, which we explore below, help inform these questions on early solar system architecture, chronology, and composition.
2.2. The chronology of the early solar system as recorded by meteorites

Meteorites provide a record of the timescales, processes, and materials that went into building the solar system (Huss and McSween, 2022). These primitive materials have a variety of physical and chemical properties that depend strongly on two factors: formation time and formation location (Ida, 2019). The formation time of the parent bodies of meteorites is a key consideration because analyses of meteorites have demonstrated that the solar system formed with significant amounts of ²⁶Al, a heat-producing short-lived radionuclide. With a half-life of roughly 720,000 years, ²⁶Al is considered the dominant heating mechanism for planetesimals (Huss et al., 2009; Lee et al., 1977). A detailed discussion of the decay scheme and heat production of ²⁶Al is provided by Castillo Rodríguez et al. (2009).

An initial abundance of 5 × 10⁻⁵ for ²⁶Al has been inferred from analyses of calcium–aluminum rich inclusions (CAIs) (Dunham et al., 2026). This would provide sufficient heat for widespread melting and differentiation (core and mantle formation) of planetesimals that formed within 2 Myr of solar system formation, depending on their composition (e.g., Al abundance and fraction of water ice). Planetesimals that form between 2 and 3 Myr may not experience metal–silicate differentiation but would be heated sufficiently to cause thermally driven recrystallization and lithification of chondritic components (Fig. 2). Temperatures for planetesimals that formed 3–4 Myr (after much of the ²⁶Al had decayed) could be sufficient to melt water ice, which would lead to water-rock alteration, and potential subsequent lithification (Castillo-Rogez and Young, 2017; Grimm and McSween, 1993; Wakita and Sekiya, 2011). However, if there is insufficient ²⁶Al to melt accreted ice, the planetesimal would remain a relatively unconsolidated sediment.

Based on composition, mineralogy, and texture, meteorites can be classified into two overall groups: those that show evidence for undergoing melting are referred to as nonchondrites or differentiated meteorites, while those that avoided melting are classified as chondrites. Chondritic meteorites are remnants of planetesimals that escaped melting and differentiation. They are aggregates of micrometer-sized particles (fine-grained dust) to centimeter-sized particles (clasts) that formed at different times and locations in the solar nebula (the forming Sun and its protoplanetary disk). Their major components are chondrules and refractory inclusions embedded in a fine-grained matrix. Chondritic components in the fine-grained matrix also include organic material and circumstellar grains (van Kooten et al., 2024). Within the solar nebula, chondrites help constrain conditions such as temperature and pressure, which varied with time, distance from the Sun, and distance from the mid-plane of the protoplanetary disk (Chambers, 2023; Ciesla and Cuzzi, 2006; Tissot et al., 2025). For instance, the reduced mineralogy and dearth of hydrated phases in enstatite chondrites are indicative of formation in a water-poor environment with elevated C/O ratios, likely in the hot inner solar nebula (within 1 au) (Ebel and Alexander, 2011; Shukolyukov and Lugmair, 2004; Wood and Hashimoto, 1993). Planetesimals that formed beyond the water ice condensation front (or snowline) accreted significant amounts of ices comprised predominately of water but also CO₂ and/or NH₃. This resulted in planetesimals with relatively oxidized compositions, such as carbonaceous chondrites (Morbidelli et al., 2016). Melting of water ice due to internal heating in planetesimals led to the formation of aqueously altered minerals in chondrites, such as hydrated silicates, carbonates, and oxides, which are abundant in CI, CM, and CR chondrites (Alexander et al., 2012; Lee et al., 2025). Several chondrite groups (OCs, CV, CK, COs, and RCs) show evidence of accretion of water ice but exhibit limited fluid alteration compared with CI, CM, and CR carbonaceous chondrites. In some cases (OCs and RCs), aqueous alteration signatures are obscured by dehydration during subsequent thermal metamorphism in planetesimals (Hutchison et al., 1987; Krot et al., 2025). Remnants of icy planetesimals occur as comets, Kuiper Belt objects (KBOs), trans-Neptunian objects, centaurs, and carbonaceous asteroids (e.g., C, B, and possibly D and P types) (Anderson et al., 2025; Demeo et al., 2022; Vernazza et al., 2021).

Geochemical analyses of these early solar system materials provide tight constraints on the timescale of solar system formation. Several radiochronometers support the idea that refractory inclusions — CAIs and ameboid olivine aggregates — are the oldest particles to have formed in the solar system (MacPherson et al., 2025; Jansen et al., 2024). Absolute ages of the oldest CAIs, dated with the ²⁰⁷Pb-²⁰⁶Pb system, indicate an age of 4567.30 ± 0.16 Myr (Connelly et al., 2012), which is also the best estimate for the formation time of the solar system. The ancient ages, refractory compositions, mineralogies, and textures of some refractory inclusions are consistent with formation via direct condensation from a gas of solar composition at temperatures and pressures expected for conditions in the solar nebula (Ebel et al., 2006). Formation of refractory inclusions during FU Orionis events has been inferred from the refractory composition, the multimodal abundances of ²⁶Al, and the occurrence of CAIs in cometary material (e.g., Boss et al., 2012; Dunham et al., 2026; Larsen et al., 2020; MacPherson et al., 2025).

Chondrules are igneous particles that make up the bulk of most chondrites. Some ²⁰⁷Pb-²⁰⁶Pb system dating analyses indicate that their formation started concurrently with CAIs (Connelly et al., 2012), although this evidence is inconsistent with relative ages determined from ²⁶Al-²⁶Mg analyses, which suggest that chondrules formed 1–4 Ma after CAIs (Nagashima et al., 2018). The mechanisms for chondrules and CAI formation are highly debated (Connolly and Jones, 2016); formation scenarios include thermal processing of dust associated with nebular shock waves (Boley et al., 2013; Desch and Connolly, 2002), polar outflows from the protosun (Liffman, 1992; Shu et al., 2001), and planetesimal collisions (Johnson et al., 2015; Stewart et al., 2025). Chondrule formation timescales and thermal models suggest that chondritic planetesimals formed 2–4 Myr after CAIs (Blackburn et al., 2017; Bouvier et al., 2007; Trieloff et al., 2003).

Differentiated meteorites include stony achondrites, irons, and stony irons, which are fragments from the rocky crusts, iron cores, and collisional mixtures of differentiated planetesimals (Chabot and Haack, 2006; Fu et al., 2017; Ruzicka et al., 2017). Iron meteorites consist mostly of Fe-Ni metal minerals called kamacite and taenite that record core crystallization and shock histories of planetesimals. ¹⁸²Hf-¹⁸²W radiometric age dating of iron meteorites indicates that they formed 1–2 Myr after CAIs at a point when a significant fraction of ²⁶Al was still extant and could facilitate widespread melting (Kruijer et al., 2014; Spitzer et al., 2021). Stony achondrite meteorites are mafic and ultramafic silicate rocks that consist of varying amounts of olivine, pyroxene, and/or feldspar minerals that record metal–silicate differentiation, and magma crystallization during crust formation on planetesimals. Hf-W ages of basaltic achondrites (angrites and eucrites) indicate metal–silicate formation occurred within 1 Myr after CAIs, while ²⁶Al-²⁶Mg chronometry indicates that magma differentiation occurred within 2–4 Myr after CAIs and crystallization occurred between 4 and 22 Myr after CAIs (Kleine and Wadhwa, 2017).

Chondritic and differentiated meteorites fall into two genetic groups based on nucleosynthetic isotopic anomalies that indicate that there were at least two chemical reservoirs, in the inner solar system and outer solar system, respectively,referred to as non-carbonaceous (NC) and carbonaceous (CC) (Warren, 2011; Lichtenberg et al., 2021). Recent studies indicate that CIs, Ryugu, and Bennu may represent another outer solar system isotopic reservoir that is distinct from the CC group (e.g., Barnes et al., 2025; Nesvorný et al., 2024; Spitzer et al., 2024). Formation of these reservoirs has been attributed to barrier gaps (optically thin orbital regions) in the disk generated by the growth of Jupiter’s core to 20 Earth masses within 1 Myr after CAIs (Kruijer et al., 2019), and the other giant planets likely experienced a similar early formation history. In contrast, ¹⁸²Hf-¹⁸²W chronometry indicates that accretion and core formation for Earth were more protracted and occurred within 30 Myr after CAIs (Fischer and McDonough, 2025; Halliday, 2014; Kleine et al., 2002; Yin et al., 2002). Moreover, the formation of the Moon occurred relatively late, approximately 120 Myr after CAIs, likely due to a collision between Earth and a Mars-sized impactor (Borg and Carlson, 2023; Yang et al., 2026). Rapid formation of planetesimals and planets and chondrule ages are consistent with disk lifetimes of a few million years based on astronomical observations of YSOs (Williams and Cieza, 2011).

2.3. Origin of the raw materials for life

The reservoirs of the six most relevant elements for life—carbon (C), hydrogen (H), oxygen (O), nitrogen (N), phosphorus (P), and sulfur (S) (referred to as CHONPS)—can be traced back to the ISM. The elements H, C, N, and O are among the most abundant elements in the universe, but abundance alone is insufficient to ensure that the element is incorporated in the chemical budget of a habitable planet. From the ISM to the initial budget protoplanetary disks, a handful of molecules are thought to be key for maintaining the abundances of C, O, and N through the processes that lead to a fully formed potentially habitable world: CO and refractory C for C, CO water and silicates for O, and HCN, ammonia (NH₃), and N-bearing organics for N (Agúndez and Wakelam, 2013; Krijt et al., 2023; Öberg and Bergin, 2021). These exact role of these reservoirs is not fully understood, particularly for O and N, and they are the subject of ongoing study (Öberg and Bergin, 2021). Other compounds relevant for the origins of life, such as methane (CH₄), HCN, methanol (CH₃OH), and many organics, can also be found in the ISM (Öberg and Bergin, 2021). Below, we review the processes that ultimately create these prebiotic reservoirs, and how they are modified until their accumulation into the planets’ building blocks. Exceptions are S and P, whose reservoirs in the ISM are not yet constrained (Öberg and Bergin, 2021 and references therein).

2.3.1. Interstellar heritage: The chemical journey from the ISM to protoplanetary disks

All stars are formed in molecular clouds, which set the initial chemical budget of protoplanetary disks. Part of the molecular cloud reservoir will be transformed into new chemical species, while other compounds will be inherited directly by the disk and remain intact after planet formation. The physical conditions (e.g., pressure, temperature, and density) required for a chemical reaction constrain the environment where a given compound is formed. Due to the low temperatures (<50 K) and densities (∼ 10⁵ cm⁻³) of molecular clouds (Chevance et al., 2020; Snow and McCall, 2006), three-body reactions occur on timescales of >10⁵ years (Agúndez and Wakelam, 2013), which means that they have no significant contribution to cloud chemistry. Gas chemistry is mostly driven by ions formed by UV photons and cosmic rays, the latter dominating the chemistry of opaque cores. Other sources of chemical compounds are grain–surface chemistry and gas grain–surface interactions.

The evidence from presolar grains found in meteorites and asteroids (Liu, 2025; Nittler and Ciesla, 2016), cometary samples (Brownlee and Joswiak, 2017; Rubin et al., 2020; Westphal et al., 2014), in situ dust measurements from the Ulysses, Cassini, and Galileo spacecrafts (Mann, 2010; Sterken et al., 2019), and astronomical observations (Draine, 2003; McClure et al., 2020; Potapov et al., 2020; Tielens, 2013) shows that dust grains are composed of silicates, graphite, silicon carbide, aluminum oxide-bearing minerals (e.g., corundum, spinel, and hibonite), hydrocarbons (aliphatic, aromatic, and hydrogenated amorphous), and ices (H₂O, CO, CO₂, NH₃, and CH₄). The most abundant ice in these mantles is H₂O, followed by CO and CO₂ (e.g., Öberg, 2016). There is still debate as to whether the silicates and the carbonaceous materials are mixed in ISM grains or are two separate populations, as well as how these populations change in different ISM environments (e.g., Compiègne et al., 2011; Hensley and Draine,2021; Jones and Ysard, 2019). Polycyclic aromatic hydrocarbon (PAH) molecules contain up to ∼20% of the elemental C (Tielens, 2008). They are formed from units made of C atoms organized in hexagonal (aromatic) rings, with peripheral H atoms. In the diffuse ISM, PAHs may contain N or deuterium (D) atoms instead of H. PAH molecules are part of the gas budget, but they are also considered grains when they form clusters with more than 10⁶ C atoms that exceed 0.01 mm (Draine, 2003).

Gas chemistry in the dense ISM is controlled by $H_{3}^{+}$ , the product of the ionization of H₂ by cosmic rays and X-rays. The $H_{3}^{+}$ radical produces CH₄ and H₂O, among other hydrogenated compounds (Öberg and Bergin, 2021). After H₂, CO is the second most abundant molecule in molecular clouds and is created by several pathways, including reactions of OH with C⁺, and from HCO⁺ (Prasad and Huntress, 1980). Nitrogen forms N₂ and NH₃; the former is the result of neutral reactions of N with OH or CH, while the latter forms from initial splitting of N₂ by He⁺, which leads to N⁺. The N⁺ ion in turn reacts with H₂ to form $N H_{4}^{+}$ , which in turn reacts with an electron (e⁻) to create NH₃ (Agúndez and Wakelam, 2013). HCN and HNC are also among the most abundant N-bearing molecules. Their formation mechanism is via dissociative recombination of HCNH⁺ with an electron (Agúndez and Wakelam, 2013). Grains act as catalysts for the formation of H₂, H₂O, CH₃OH, and more complex molecules through the addition of H and/or recombination of radicals that are produced by UV or electro-mediated dissociation (Öberg and Bergin, 2021).

From the outermost regions of a molecular cloud to the center of a pre-stellar core (Fig. 1), the temperature decreases while the density increases, which also increases the opacity for UV photons, and these temperature and density gradients in the parent cloud have consequences for the chemistry of the disk. At the edge of the cloud, where the temperature is higher, H₂ is formed on the surface of bare silicate grains, and CO is produced in the gas phase. Deeper in the cloud, H₂O is created and freezes on the grain surface. In the pre-stellar core, the temperature is low enough (<20 K) for CO to freeze out over the grains, which allows N₂H⁺ to accumulate and produce N₂ (Öberg and Bergin, 2021). The center of the core is shielded from UV photons, and the temperature is low enough (∼10 K) for more volatile species (CO₂, NH₃, CH₄, and N₂) to condense onto grains. Ices in the pre-stellar core are part of the initial volatile inventory available for planet formation (Öberg and Bergin, 2021).

From the ∼300 molecules identified in the ISM, about 50% of them are interstellar complex molecules (COMs or iCOMs) that contain six or more atoms where at least one of them is C¹. Experiments with ices irradiated by photons or energetic particles have shown that most relevant prebiotic compounds, including amino acids, sugars, and nucleobases, can be formed in icy mantles (Sandford et al., 2020). The products of ice irradiation are similar regardless of the source of energy (photons or particles) or the sources of C, H, O, and N (Sandford et al., 2020 and references therein). Experiments with ices irradiated with UV found that the relative abundances between amino acids are similar to those observed in meteorites. This suggests a link between pre-stellar cores, the building blocks of planets, and prebiotic chemistry (Sandford et al., 2020). In the coldest regions (∼10 K) of pre-stellar cores, the sublimation of grains’ icy mantles cannot be the main source of the COMs detected in the gas phase (Ceccarelli et al., 2023 and references therein), which implies that gas chemistry also contributes to the production of organics such as HCN and unsaturated (low H/C ratios) carbon chain molecules (Agúndez and Wakelam, 2013; Ceccarelli et al., 2023; Öberg and Bergin, 2021).

After the pre-stellar core infall begins, radiation from the protostar enables new chemical pathways. At the outer region of the Class 0 object, the cold (10–30 K) envelope preserves the ice mantles of the grains. At above 20 K, radicals can diffuse in the ice, producing organics; at hotter (30 K) regimes, CO, N₂, and CH₄ sublimate, and CH₄ in the gas phase produces unsaturated carbon chains (cyanopolyynes). Water and organics are released to the gas phase at the zones where the envelope temperature exceeds 100 K and where the mild shocks are produced by infalling envelope material on the disk surface. Disk densities allow gas chemistry in the inner disk (<10 au) that may include three-body reactions (Öberg and Bergin, 2021). Class 0/I objects are characterized by outflows that do not have a relevant role in disk chemistry but serve as probes of disk and envelope chemical composition. For example, the shock of the outflow with the envelope contains organic molecules that may result from the rapid sublimation of ice mantles of grains in the envelope (Ceccarelli et al., 2023; Öberg and Bergin, 2021; van Dishoeck, 2003). The combination of observations with ALMA and the Herschel Space Observatory, along with developments in chemical modeling and laboratory experiments, has shown that Class 0/I objects have a rich organic chemistry that results from the combined contribution of warm (∼20 K) and cold grains and gas-phase chemistry (Jørgensen et al., 2020). Once the envelope is cleared, the disk temperature is set by radiation from the protostar; at this stage, the YSO is a Class II object. In the context of the solar system, this is considered “time zero” and defined in meteoritics as the condensation of CAIs, ∼ 4.567 Gyr ago (Fig. 2; Connelly et al., 2012).

Discussion is ongoing regarding the chemical pathways within disks and the influence of the parent molecular cloud, including understanding in situ versus inherited disk chemistry (e.g., Öberg et al., 2023 and references therein). State-of-the-art astronomical facilities have significantly advanced our understanding of disks, given the progress of high-resolution instrumentation leading to new interpretations of YSO pathways (e.g., Brittain et al., 2005; Goto et al., 2003; Smith et al., 2015). The NASA Infrared Telescope Facility (IRTF) has recently been used to evaluate O reservoirs toward massive YSOs in the solar neighborhood, tracing gas that is largely the outer envelope and parent molecular cloud. Ratios of ¹⁸O/¹⁶O and ¹⁷O/¹⁶O, derived from CO, fall along a mass-independent fractionation line in three-isotope space with respect to the local ISM and are evidence of CO self-shielding in predominantly cloud (i.e., non-disk) material. These observations suggest that CO self-shielding could be inherited from the parent cloud reservoir (Smith et al., 2024) rather than a disk phenomenon (e.g., Öberg et al., 2023). We explore this further in the following section.

2.3.2. Chemical composition of the early solar system inferred from disk observations and meteorites

Disks of Class II objects have been extensively studied using observations, laboratory experiments, and numerical models. As discussed above, the initial chemistry of the disk is a mixture of materials inherited from the molecular cloud, compounds processed during the collapse, and products of disk chemistry (Öberg and Bergin, 2021 and references therein). When the disk is formed (Class 0/I stage), the infalling material close to the protostar is heated, which breaks molecules into atoms that recondense to form new chemical species. The disk spreads, and these new materials are transported to larger disk radii dominated by pristine materials from the molecular cloud or the disk envelope. Models and solar system D/H water ratios suggest that water in protoplanetary disks is mostly inherited from the ISM (references in Öberg et al., 2023). Silicates, in contrast, must be processed at the hot temperatures (>1000 K) within the disk to reach the high crystalline fraction relative to amorphous silicates (∼20%) systematically observed in IR compared with the ISM silicates (<2%) (Natta et al., 2007; Oliveira et al., 2011). Although, depending on the geometry of the system, IR observations may sample only the surface of protoplanetary disks. The model from Jang et al. (2024) predicts that due to the short vertical mixing timescale for the small grains (<3 µm), the abundances of crystalline silicates derived from IR observations are representative of the disk mid-plane.

Protoplanetary disks are typically divided radially into two sections, inner and outer disk. The division between both sections is set around 10 au from the star (e.g., Henning and Semenov, 2013; Öberg et al., 2023), but for specific studies of the solar system, the inner disk is typically referred to the formation region of the rocky planets (<5 au) (e.g., Bizzarro et al., 2025; Broadley et al., 2022). In all cases, the inner disk is characterized by its temperatures; far- and mid-IR observations of Class II protoplanetary disks indicate that mid-plane temperatures close to the star (<3 au) range from 700 to 1700 K ( $\leq$ 0.2 au) to 500 K at around 1 au (Miotello et al., 2023). Molecules detected in the IR from this region are OH, CO, H₂O, CO₂, and HCN (references in Miotello et al., 2023; Öberg and Bergin, 2021). The water snowline is located where the temperature is ∼150 K, low enough to condense water onto grains; consequently, this molecule is not available for gas-phase chemistry. The position of the snowline depends on the temperature profile of the disk, which is mostly determined by the protostar flux and accretion rate. Water condensed into grains at and beyond the snowline can potentially be accreted into planetesimals and later onto protoplanets. Meteorites show a more complex scenario for the solar system, where planetesimals can lose their volatiles due to internal heating even if they form beyond the snowline (Grewal et al., 2024; Newcombe et al., 2023).

Beyond the water snowline is the outer disk, which is stratified into three layers related to chemistry: (1) The first layer is the disk surface, where the chemistry is dominated by radiation. This layer is typically referred to as the photon-dominated region because stellar UV and cosmic rays drive the chemistry by ionizing and photodissociating molecules. Here, X-rays ionize He, which produces He⁺ that destroys CO molecules and thereby initiates a rich organic chemistry (Henning and Semenov, 2013; Raul et al., 2025); (2) The middle layer, between the surface and mid-plane of the disk, is often considered the “molecular layer” because the temperature range (30–70 K) allows gas-phase reactions; (3) At the third layer, in the disk mid-plane, the gas is mostly protected from high-energy radiation, and temperatures <20 K result in the freezing of molecules over the grains. In this third layer, sources of ionized atoms and molecules are likely cosmic rays and the radiogenic decay of short-lived radionuclides (Eistrup et al., 2016; Pacetti et al., 2022).

Carbon monoxide has been observed at very high spectral resolution in a few disks at particular geometries that allow for the evaluation of disk gas along an integrated line of sight. Such observations are significant toward understanding the origin of solar system isotope anomalies, with implications for the inheritance of key chemical pathways in potential new Earth-like systems. With the use of VLT-CRIRES to obtain near-IR absorption spectra of CO in rovibration (2 and 4 mm), mass-independent fractionation of O has been observed, which lends critical observational support for CO self-shielding (Smith et al., 2009) as an explanation for the solar system O isotope anomaly, a decades-long mystery discovered in primitive meteorites (Clayton et al., 1973). It is unclear whether CO self-shielding is a disk phenomenon or rather is inherited from the parent cloud; competing scenarios that dictate cloud-disk evolutionary pathways that lead to planets and prebiotic molecules. However, recent high-resolution observations with the iSHELL spectrograph at the NASA IRTF suggest that CO self-shielding might be inherited from the parent cloud in YSOs within 1 kpc of the Sun (Smith et al., 2024). Additionally, observations that trace outer envelope/parent cloud toward several massive YSOs beyond the solar neighborhood thus far show significant ¹⁶O-excesses relative to ¹⁷O and ¹⁸O in comparison to predictions from galactic chemical evolution models (e.g., Prantzos et al., 1996), which suggests that inheritance of molecules from CO self-shielding chemical in the cloud may be a wider galactic phenomenon (Smith et al., 2025). Continued work with state-of-the-art observatories will build these data sets and add significance to these findings. With respect to carbon isotopic reservoirs, analyses of disks and envelopes observed with VLT-CRIRES have revealed heterogeneity in [¹²CO]/[¹³CO] with values significantly higher than local ISM (∼68) and solar (∼87), with suggested interplay between CO ice and gas reservoirs influencing gas-phase C isotope ratios (Smith et al., 2015). Laboratory experiments with interstellar ice analogs have further shown that such heterogeneity could be the result of ¹³CO being more tightly bound in the ice phase (Smith et al., 2021). Such observations are significant toward understanding the origin of solar system O isotope anomalies and key chemical pathways toward potential new Earth-like systems.

Complex organic molecules are challenging to detect or identify in disk spectra even with the most powerful instruments, such as ALMA (Öberg et al., 2021) and JWST (van Dishoeck et al., 2023). From the total budget of organic molecules in the disk, we can only observe the molecules in the gas reservoir. Models suggest that most organic molecules in the disk are locked in the ice reservoir, where they cannot be detected (Öberg et al., 2021). Disk organic chemistry is characterized by fewer O-bearing compounds, compared with molecular clouds and pre-stellar cores (Öberg et al., 2023; Zhang, 2024). For example, CH₃OH is readily detected in pre-stellar cores but has only been detected in one T-Tauri star (Walsh et al., 2016). Within a few astronomical units from the star (inner disk), the organics detected (e.g., C₂H₂, HCN, and CH₄) are consistent with high-temperature gas chemistry (Öberg et al., 2023 and references therein). In the outer disk, detected complex organics include C₂H, H₂CO, c-C₃H₂, and nitriles (e.g., HC₃N and CH₃CN) (Öberg et al., 2023,2021). Using the gas reservoir budget of organic molecules measured by ALMA (Table 1, Drozdovskaya et al., 2019), recent models predict a minimum organic reservoir equivalent in mass to 2–460 Earth oceans within 50 au in five protoplanetary disks (Table 8, Öberg et al., 2021). The calculated organic budget is <3.5% of the water abundance in the studied disk sample (Öberg et al., 2021). Analogous to the snowline, a tar or soot line has been proposed as the disk radius where the temperature is ∼500 K and refractory organics are mostly lost by sublimation (Gail and Trieloff, 2017; Li et al., 2021; Lodders, 2004). Models predict that this line would move toward the Sun and fall closer than 1 au within 1 Myr. The loss of solid C-bearing species would result in planetesimals with lower volatile content formed at Earth’s accretion zone (Li et al., 2021). Other possible scenarios for C depletion for the building blocks of our planet are the loss of C-bearing species by internal heating of planetesimals (e.g., Bergin et al., 2015; Hirschmann et al., 2021) and the dominance of refractory inclusions and chondrites already depleted of volatiles in Earth’s building blocks (Alexander, 2022).

Dust grains grow and settle to the disk mid-plane. During this process, volatiles condense over the grains, which removes these compounds from the available budget for gas chemistry (Öberg and Bergin, 2021). The formation of pebbles (see Section 2.1.1) decreases the opacity of the disk and thus modifies the temperature and the radiation shield; this stage of the disk is not completely understood, but observations suggest that there will be an accumulation of materials at the major volatiles’ snowlines (H₂O, CO, and CO₂) and the formation of pressure bumps that act as traps for solids and volatiles (Öberg and Bergin, 2021). Observations of C, N, and O abundances relative to H in the TW Hya disk (T-Tauri ringed disk) suggest that the volatile budget retained in solids in the gas traps is similar to that of CI chondrites (McClure et al., 2020).

Chondrites have bulk compositions that are strikingly similar to the solar photosphere for all elements except for the most volatile elements, the noble gases, H, C, O, and N, all of which condensed as ices and icy planetesimals in the outer solar system (Aikawa et al., 1999; Lodders, 2021). There is a limited meteoritic record in chondrites, because ices would have been readily lost via sublimation and/or evaporation as these objects enter the inner solar system (within 3 au). Carbonaceous chondrites such as CI, CM, and CR chondrites show evidence of substantial fluid alteration, and many lines of evidence suggest that these fluids were derived from H₂O ice that melted due to the heat from ²⁶Al decay (Lee et al., 1977,2025). This alteration resulted in abundant hydrated silicates in the matrix of the most pristine chondrites, including ordinary chondrites (Chan and Zolensky, 2022). Carbonate minerals (e.g., calcite [CaCO₃], etc.) also occur in aqueously altered chondrites, which indicates that the altering fluids and/or ices contained CO₂ or CO (Alexander et al., 2015). Ammonia ices have also been inferred for carbonaceous chondrites based on direct detection in Bennu samples (Glavin et al., 2025) and indirectly from spectral analyses of hydrated asteroids (Rivkin et al., 2022). Organic compounds also occur in the fine-grained matrix of pristine chondrites as mostly insoluble kerogen-like PAHs (up to 3.5 wt.%), but compounds highly significant to astrobiology, such as amino acids, sugars, and purines, have also been identified at abundances up to 600 ppm (Glavin et al., 2018; Martins and Sephton, 2009). These soluble organic compounds can form molecules, such as carbohydrates, lipids, proteins, and nucleic acids, that are essential for the development of life on Earth (Chan and Zolensky, 2022). The Primordial Multifunctional Organic Entities–PriMEs scenario proposes that insoluble organic matter in carbonaceous chondrites and interplanetary dust played a relevant role in the origins of life (d’Ischia et al., 2021). This model is based upon the structural and physicochemical properties that are shared by insoluble organic matter in carbonaceous chondrites and melanin polymers present in bacteria and fungi (d’Ischia et al., 2021).

Carbonaceous chondrite material returned from hydrated asteroids Ryugu (C-type) and Bennu (B-type) contains similar varieties and abundances of organics. This lends further support to the idea that extraterrestrial materials would have delivered prebiotic molecules to Earth (Glavin et al., 2025; Naraoka et al., 2023; Parker et al., 2023). In addition to the composition of ice and organics inferred from chondrite analyses, spectral surveys of comets, KBOs, and other remnants of icy planetesimals indicate that refractory C, CO₂, CO, CH₄, NH₃, HCN, and simple alcohols such as ethane and methanol also occur in ice-rich planetesimals (Bockelée-Morvan and Biver, 2017; Brown, 2012; Schaller and Brown, 2007). The relationship between hydrated asteroids and KBOs is inferred by the spectral similarities between these objects, which suggest large-scale transport of dust and planetesimals between the inner and outer solar system (Campins et al., 2010). Recent analyses of pristine samples from Bennu have found high abundances of ammonia and enrichments in N isotopes, consistent with formation in the outer solar system (Glavin et al., 2025). Dynamical mixing of planetesimals between the outer and inner solar system (where Bennu’s parent body formed vs. where it is currently located) is thought to have played an important role in the delivery of water, organics, and other life-building volatile ingredients to the region of Earth’s formation (Alexander et al., 2017; Piani et al., 2020; Steller et al., 2022).

In the solar system, we can probe the earliest stages of planet formation and our interstellar heritage through the composition of comets, asteroids, and meteorites, which are relics of the planetesimals that build our habitable planet. The following section further addresses these building blocks.

2.3.3. The planetary building blocks: Link between the ISM and planets

Planetesimals, the building blocks of planets, are formed from dust when the meter-size barrier is overcome (see Section 2.1.1). The most promising hypothesis for planetesimal formation is the initial formation of 50–100 km sized bodies by streaming instabilities. Theoretical predictions from gravitational instability models closely agree with the inflection in the size distribution of asteroids (Krijt et al., 2023, and references therein). Moreover, the characteristics of the dust aggregates and porosities in comet 67P/Churyumov-Gerasimenko (67P/C-G) and the rubble-pile asteroid Ryugu are consistent with pebbles accreted in a gentle gravitational collapse (Watanabe et al., 2023 and references therein). Comets and asteroids can be considered fossil planetesimals for a few reasons: they were accreted in the protoplanetary disk that may not have become internally differentiated, they are fragments of collisional events, or they were accumulated from debris of previously disrupted planetesimals. Regardless of their origin, comets and asteroids are integral to the understanding of planet formation.

Comets and asteroids were once regarded as two chemically separate reservoirs, but recent evidence suggests that there is a continuum between these groups (Bergin et al., 2024). The transitional objects between comets and asteroids would form beyond the water snowline, where they would have been dynamically perturbed during the formation and migration of giant planets (Bergin et al., 2024). This would further imply that low-albedo asteroids and carbonaceous meteorites may have come from the cometary formation zone beyond the water snowline (Bergin et al., 2024 and references therein; Bizzarro et al., 2025). Comets contain pristine materials from the ISM and early stages of stellar formation, shown via comparisons of the organics detected in comets with pre-stellar cores and molecular clouds (Altwegg et al., 2019). For example, the abundances of CHO-bearing molecules relative to CH₃OH of 67P/C-G, as measured by the Rosetta mission, are similar to the abundances observed by ALMA in the IRAS 16293-2422 B disk (Drozdovskaya et al., 2019). In contrast, cometary material collected from the coma of the Jupiter Family Comet Wild 2 by the Stardust mission contained CAIs and chondrules that were processed at high temperatures (see Section 2.2), and mineralogies similar to CAIs and chondrules in carbonaceous chondrites (Sandford et al., 2021). Mixtures of pristine ISM materials and disk-processed components are also observed in asteroids. The OSIRIS-Rex mission returned samples of asteroid (101955) Bennu, which reveal a mixture of inherited materials, including presolar grains (Lauretta et al., 2024), and a mineralogy that suggests processing in the protoplanetary disk (Russell et al., 2024). The combination of pristine materials and components processed at high temperatures is evidence of radial mixing in the protoplanetary disk that formed the solar system (Altwegg et al., 2019 and references therein).

The favored current hypothesis concerning the origin of Earth’s water is that it was accreted during the final stage of the planet’s formation, from bodies formed beyond the water snowline with compositions analogous to carbonaceous chondrites (Meech and Raymond, 2020). This is further supported by ¹⁵N/¹⁴N, D/H, and elemental ratios of noble gases that point to a contribution of 30–60% of carbonaceous chondrite type materials as the main source of volatiles during the final stages of Earth’s accretion (Broadley et al., 2022 and references therein).

3. The Early Sun

3.1. Bolometric luminosity of the early Sun and its evolution

Main sequence stars fuse H into He in their cores; this generates energy that creates pressure to halt their gravitational collapse (Christensen-Dalsgaard, 2021; Salaris and Cassisi, 2005). This energy ultimately escapes from the stellar photosphere and can be intercepted by orbiting planets to warm their atmospheres and surfaces. A star’s luminosity evolution is controlled primarily by its mass, with second-order effects from metallicity (the abundance of elements other than H and He). Before the onset of fusion in the core, a star’s luminosity is instead controlled by the release of gravitational potential energy during its contraction phase (Hayashi, 1961; Stahler and Palla, 2004). The star can be substantially brighter during this time than when it enters the main sequence (Kim et al., 2002; Yi et al., 2003). This phase would have lasted 30–50 million years for the Sun, which has implications for the climatic state of Earth’s earliest atmosphere after the completion of planet formation (Zahnle et al., 1988).

The luminosity of main sequence stars can be directly related to mass by the power law relation $L \propto M^{n}$ , where L is the luminosity, M is the mass, and n ranges between 2 and 5 depending on the mass range (Salaris and Cassisi, 2005). For a star of similar mass to the Sun, n is often given as ∼4.5. The time-dependent evolution of the Sun is well understood. Standard solar models have been developed through detailed physical modeling coupled with validation by observational helioseismology, the characterization of solar neutrinos, and observations of stellar clusters of known age (Bahcall et al., 2001; Christensen-Dalsgaard, 2021; Gough, 1981; Ribas et al., 2010). From these models, we can infer that the early Sun was ∼70% as luminous as it is today when it entered the main sequence and was 20–25% less luminous during the Archean Eon (4.0–2.5 Ga). The Sun brightens as H in the core fuses into He, transforming four protons into one He nucleus and thus decreasing the core’s mean number density. In response, the core contracts, which increases temperature and density to counter the force of gravity. The rate of fusion is highly dependent on temperature, so as the core contracts, the energy generation rate increases, pushing out the outer layers and increasing the radius of the Sun. This increase in radius is the primary driver of the Sun’s luminosity increase, rather than changes in the effective temperature of the solar photosphere (Bahcall et al., 2001). Figure 3 illustrates the reconstructed luminosity evolution of the Sun based on stellar models (Kim et al., 2002; Ribas et al., 2010; Yi et al., 2003).

FIG. 3.

Time evolution of the luminosity, radius, and temperature of the Sun in comparison to current values. Figure by R J Hall CC BY-SA 3.0 Wikipedia Commons, based on Figure 1 in Ribas et al. (2010).

3.1.1. Influence on early climate and the Faint Young Sun problem

Sagan and Mullen (1972) famously described the “Faint Young Sun” problem (FYSP) based on coupled solar and climatic modeling, which shows that early Earth would have been frozen with modern atmospheric composition; this contrasts with geological evidence that early Earth was not globally glaciated. Yet, earlier authors identified a tension between a predicted faint early Sun and habitable early Earth (Donn et al., 1965; Schwarzschild et al., 1957). Various solutions have been proposed for the FYSP, including additional greenhouse gases, reduced cloud cover, and greater fractional ocean coverage (see reviews in Feulner, 2012; Spencer, 2019). Much of this climate problem is likely resolved by higher levels of CO₂ on early Earth (Charnay et al., 2020), maintained by balances between volcanic outgassing and geochemical weathering processes (Krissansen-Totton and Catling, 2020; Walker et al., 1981). It is even possible that the reduced luminosity of the young Sun facilitated water condensation and thus habitable conditions on early Earth’s surface, as water clouds on the nightside generate a substantial greenhouse warming that may have forestalled the formation of liquid oceans on the more irradiated Venus (Turbet et al., 2021). We do not discuss possible planetary solutions to the FYSP at length here, as the evolution of Earth’s atmosphere and climate—including consideration of the FYSP—is covered more extensively in other works (e.g., Spencer, 2019; Catling and Zahnle, 2020; Charnay et al., 2020).

Here, we briefly address and dismiss one proposed solution for the FYSP that involves the characteristics of the early Sun. Because solar luminosity is tightly controlled by mass, some authors have suggested that a more massive early Sun could have maintained early Earth’s temperate climate (e.g., Graedel et al., 1991). Because the orbit of Earth would also have been controlled by solar mass (Minton and Malhotra, 2007), the insolation received (F) would have an exceptionally strong dependence on mass ( $F \propto M^{6.5}$ ). While an early Sun only ∼6% more massive would be sufficient to resolve the FYSP, it would require average mass loss rates several hundred times greater than today, inconsistent with those inferred for younger Sun-like stars (Spencer, 2019). Recent solar evolution modeling constrained by observed mass loss rates suggests largely insignificant deviations from past predictions and a minimal role for a more massive early Sun in resolving the FYSP (Basinger et al., 2024).

3.1.2. Stellar evolution and habitable zones

The “habitable zone” is defined as the range of distances around a star where an orbiting rocky planet could maintain surface liquid water, depending on its atmospheric composition (Kasting et al., 1993; Kopparapu et al., 2013). As stars increase in luminosity as they age, the habitable zone boundaries also move outward. More massive stars have shorter lifetimes and increase in luminosity more quickly, while less massive stars live longer and increase in luminosity more slowly. The range of distances where a rocky planet can sustain liquid water on its surface for a given period, such as from the onset of H fusion in the host star to the current epoch, is called the “continuously habitable zone.” This concept may be important for exoplanetary systems because some planets will have formed outside the habitable zone and moved into it as their star brightened (Tuchow and Wright, 2020). Such worlds, unlike Earth, may have begun their geologic history fully glaciated, a so-called “cold start” (Kasting et al., 1993), which may complicate later habitability (Wolf et al., 2017). Regardless of whether a potentially habitable planet is in the continuously habitable zone, it is important to consider that any planet’s climate would have been influenced by the evolving bolometric luminosity of its host star over geologic time.

3.2. The active young Sun

Astronomers probe the conditions likely present on the early Sun through the use of solar analogs of differing ages. Conditions present on the very early Sun (within roughly 10 million years of formation) are complicated by the fact that emissions originate from both the star itself and mass transfer from the remnants of a planet-forming disk. This latter accretion phenomenon reveals itself as a ∼10,000 K optically thick plasma in the UV spectral region (Johns-Krull et al., 2000). Young stars in this age range without disks also exhibit enhanced amounts of UV emission, attributable to the increased levels of chromospheric activity due to the stars’ fast rotation by way of its youth. The well-known rotation–activity relation (Skumanich, 1972) is a fundamental connection between magnetic field generation via a dynamo process that leads to the observed “activity” and the stellar rotation. Stellar mass loss from a magnetized stellar wind sheds angular momentum and acts to slow the star down with time. At these young ages, the stars are rotating quickly and hence have correspondingly high levels of X-ray and UV emission from magnetic activity tracers (starspots, flares, and chromospheric and coronal emission); these emissions are characterised as follows:

X-rays (<10 nm): During the first ∼5 Myr, 1 M_⊙ stars emit constant X-ray luminosity (L_X) probably from their extended corona, which is insensitive to the stellar radius changes during this first stage (Class II) (Getman et al., 2022). Young solar-mass protostars have fully convective interiors that become partially radiative, which results in a transition of the magnetic field generation mechanism that produces a power law decay in L_X with time (Getman et al., 2022; Preibisch and Feigelson, 2005). This implies that the early Sun’s L_X was around 10³⁰ erg s⁻¹, while its present-day luminosity ranges between 10²⁶ and 10²⁷ erg s⁻¹ (Getman et al., 2022; Glassgold et al., 2005). The shape of the solar spectrum at a younger age in this wavelength range was also harder with more photons in the 1–10 keV energy range.

Extreme and far-UV (l < 200 nm): A solar-like star does not settle onto the main sequence until it is roughly 40 Myr in age (Kunitomo et al., 2017). At this point, its evolution is governed largely by the nuclear forces at work in the stellar interior, as well as magnetic forces largely controlled by stellar rotation and manifesting as activity signatures. Enhanced UV emission from the star’s chromosphere is one such activity signature. Guinan et al. (2003) studied the integrated flux in the 92–118 nm band (Fig. 4, top right) for a sample of solar-like stars and noted that the solar far-ultraviolet (FUV) emissions were about twice the present value 2.5 Gyr ago and about 4 times the present value 3.5 Gyr ago. They studied individual solar-like stars with known young ages: EK Dra (130 Myr), π¹ Uma (300 Myr), κ¹ Cet (650 Myr), β Com (1.6 Gyr), as well as solar-like stars older than the Sun (β Hyi [6.7 Gyr] and 16 Cyg A [9 Gyr]). They also noted that the FUV flux of the zero-age main-sequence Sun could have been as much as 50 times higher. This study established power law relationships between stellar surface flux and rotation period; it determined values of y in the relationship F ∝ P_rot ^y . The values of the various chromospheric indicators were in general agreement, while the decay of the one coronal line in the study (Fe XVIII) exhibited a much steeper decay. Guinan and Engle (2007) (Fig. 4, bottom left and right) show the 103.0–104.0 nm and 125.0–170.0 nm spectral regions for most of these stars.

FIG. 4.

Top: Evolution of different FUV emission line and continuum markers in Sun-like stars of different ages (Guinan et al., 2003). The negative numbers to the right of each set of lines indicate the power law relationship (y in the relationship F ∝ P_rot^y). Lower left: Evolution of 103.0–104.0 nm spectral region. Lower right: Evolution from 125.0 to 170.0 nm for Sun-like stars at a range of ages from 100 Myr to 6.7 Gyr, from Guinan and Engle (2007). FUV, far-ultraviolet.

Near-UV ( $λ$ > 200 nm): The UV emission of the young Sun is considered well-known, thanks to observations of young Sun-like stars that provide proxy estimates for solar UV emission through time (e.g., Ribas, 2009; Ribas et al., 2005). These observations have been used to calibrate solar flux models to obtain estimates of solar UV emission as a function of age. Most commonly used is the publicly available model from Claire et al. (2012), which is valid from 0.1 nm to 160 μm, and from 4.4 Ga to 5.6 Gyr. Scrutiny of these models (Claire et al., 2012) has revealed a prebiotic stellar UV SED (4.4–3.7 Ga; Ranjan and Sasselov, 2016; Pearce et al., 2018) not too dissimilar from the modern solar UV SED. The early Sun emitted less UV than the modern Sun at NUV wavelengths (≥200 nm) but more UV light than the modern Sun at FUV wavelengths (Claire et al., 2012). This is because solar NUV is partially formed in the photosphere, which was cooler on the young Sun, while FUV emission is sourced from the chromosphere, which was more active, due to the young Sun’s stronger rotation. The differences between the SEDs are strongest at XUV (X-ray + EUV, l < 110 nm) wavelengths; at UV wavelengths, the band-integrated UV emission of the 4.4 Ga Sun varies by <40% relative to modern solar band-integrated UV.

Flare activity: Flaring variability occurs across the electromagnetic spectrum, and young solar-analog stars frequently exhibit flares, as expected due to their youth. These flares are also more energetic than those that occur during the main sequence. Ayres (2015) reported on a large UV flare seen on the young solar analog EK Draconis (50 Myr, G1.5V), which was a long-duration flare, with a flare decay timescale of ∼7000 s. Spectroscopy further revealed evidence for redshifts during the flare, indicating that material is falling back to the stellar surface. Electron densities were found to be about 10¹¹ cm⁻³, consistent with densities derived from large stellar coronal flares. This event was unusual in being detected with time-resolved spectroscopy; Loyd and France (2014) found no evidence for UV flares on the handful of archival observations of other young solar analogs. Brasseur et al. (2019) found photometric evidence for NUV flaring from predominantly solar-like stars observed with both the GALEX and Kepler spacecraft. Stellar ages are not known for this stellar sample; however, there is a general trend for the more energetic flares to occur on stars with short rotation periods, an indirect indicator of youth. This is also confirmed by the positive correlation between the quiescent stellar NUV luminosity and flare energy, as more magnetically active young Suns would have enhanced levels of flaring and non-flaring emission. These flares tend to be fairly short duration, lasting <300 s, with enhancements above the quiescent flux level ranging from 1.5 to 1700, and energies from 10³⁴ to 10³⁷ erg (Brasseur et al., 2019, Fig. 17), well above the largest flare energies observed on the present Sun.

Super-flares (E_X > 10³⁴ erg) and mega-flares (E_X > 10^36.2 erg) are observed in YSOs from Class I to Class III, although their rates decrease with age (Getman and Feigelson, 2021). During their first 5 Myr, solar-mass PMS stars produce roughly 10⁷ mega-flares that contribute from 10% to 20% to the total X-ray energy budget (Getman and Feigelson, 2021). Extrapolating from these observations, we can infer that PMS solar-mass stars produce 10⁶ or more super-flares than their older counterparts already in the main sequence. Neither a mega- nor a super-flare has been observed in the Sun, yet there is evidence of this activity in the early solar system. Wolf-Rayet winds from massive stars may be the source of short-lived radionuclides (T_1/2 $\leq$ 5 Myr) for the solar system (Desch and Miret-Roig, 2024), although the abundance of ¹⁰Be/⁹Be in CAIs may be explained either by galactic cosmic rays accelerated by supernova shocks (Dunham et al., 2022) or particles accelerated during flares (Getman and Feigelson, 2021; Gounelle et al., 2006).

3.2.1. The early Sun activity, its protoplanetary disk, and early Earth

The chemical structure of the protoplanetary disks is the consequence of the effects of stellar radiation on the disk gas and dust (see Section 2.1). High-energy photons and particles (X-rays, UV, and cosmic rays) ionize and photodissociate molecules; this drives chemistry in the protoplanetary disk (see Section 2.3). They are also responsible for photodesorption, a process where molecules are desorbed from ices by photons, enriching the gas phase with water and organics (Basalgète et al., 2021; Potapov et al., 2019). Numerical models show that X-ray flares affect the abundance of O-, S-, and C-bearing species, which suggests that flares may be relevant to the formation of complex organic molecules (Waggoner and Cleeves, 2022).

The lifetime of a protoplanetary disk is mostly determined by the competition of the gas accretion rate toward the star and the gas loss via photoevaporation driven by high-energy radiation that effectively warms the disk gas (Ercolano and Pascucci, 2017). Both processes impact planet formation. While there is gas in the disk, planets migrate, and their orbits move to resonant chains; once the gas is lost, planetary systems become unstable (Raymond et al., 2020). Thus, the timing of gas loss has consequences for the final structure of a planetary system. For the early solar system, there is still debate about the mechanisms that determined the masses, compositions, and positions of the inner planets and the Asteroid Belt (e.g., Raymond et al., 2020). One of the key ingredients in the predominant models for the origin and evolution of the inner solar system is the moment at which the disk gas was lost (Raymond et al., 2020). Enhanced early solar activity could have contributed to the gas loss from the solar protoplanetary disk and thereby could have influenced the volatile budget throughout Earth’s formation, which extended from the pebble-formation era (Fig. 2) to more than 100 Myr.

Once our planet was formed, its atmosphere was initially highly reduced (H-rich) and then was replaced with a neutral H₂O-CO₂-N₂ atmosphere (Zahnle et al., 2010). The duration of the period of the primordial atmosphere could be relevant to the first steps toward the origin of life. The classical Miller–Urey experiment, which attempted to reproduce Earth’s early conditions under the assumption of a reduced atmosphere, resulted in a rich mixture of organic compounds that included those that build life (Miller, 1953). Experiments with less reduced atmospheres have lower yields of compounds such as HCN, nitriles, and amines, while they also form more carboxylic acids and aldehydes (Trainer, 2013). The duration of the reduced atmosphere depends on the many unconstrained variables, such as the lifetime of the magma ocean and the atmospheric pressure that could be as large as thousands of bar (Zahnle et al., 2010). Another relevant variable is the solar activity that produced particles that could contribute to the formation of prebiotic relevant compounds (e.g., Miyakawa et al., 2002), while XUV may have contributed to the loss of the primordial atmosphere (Getman et al., 2022).

4. Summary

This article presents a current review of early solar system history. It traces key events and processes as evidenced from meteorites and astronomical observations. While some gaps may remain, we have aimed to present the latest relevant lines of evidence that should be considered in ongoing studies of the habitability of Earth-like worlds. We cover the evolutionary sequence of YSOs in the path from star formation to planetary system, from the ISM to the molecular cloud and embedded core, through the Class III stage, to the main sequence star and grain growth to differentiated planets and the origin of the primordial atmosphere.

Primitive material from planetesimals is preserved in the Asteroid Belt, Kuiper Belt, and other reservoirs of small bodies, and meteorites and samples returned from these reservoirs enable the study of timescales and processes that went into building the solar system. Primitive chondritic material, with compositions similar to the solar photosphere, contains a host of nebular particles that record nebular conditions and processes, while differentiated meteorites are the records of planetary differentiation. Isotopic analyses of chondritic and differentiated meteorites indicate that there were two chemical reservoirs, NC and CC, which source inner and outer solar system planetesimals, respectively. Radiometric analyses of meteorites indicate that within the first few million years of solar system formation, chondritic components formed and planetesimals differentiated. Evidence from meteorites indicates that carbonaceous chondrites such as CI, CM, and CRs show substantial alteration from fluid derived from melting of ice, including H₂O, CO₂, and NH₃, with large-scale mixing of dust and planetesimals between the inner and outer solar system reservoirs.

From the astronomical perspective, we emphasize recent work that shows the critical role that observations of protoplanetary disks as analogs to the solar nebula play in providing windows into planet-building and prebiotic chemical pathways. Theoretical models coupled with data from state-of-the-art astronomical observatories have enabled recent insights into the disk phases of planet formation, from in situ disk chemistry to scenarios of inheritance from the parent cloud to the disk. Disk snowlines set the limit between the gas and condensed reservoir of a given compound; from the star to the outer edge of the disk, the most relevant snowlines are for refractory carbon, water, CO₂, CO, and N₂, all of which are relevant molecules for planet formation. Carbon is the building block of all life as we know it, and understanding its chemical pathways in forming Earth-like systems inevitably requires going back to the original reservoirs. Carbon in the ISM is in the form of gas (CO, CO₂, and small organics) and refractory inorganic and organic dust. Part of the carbon budget is preserved far from the star in objects such as comets and carbonaceous meteorites, which yield the primordial source of volatiles available during planet formation. For the forming Earth, volatile inventories were likely scarce, acquired only in the final stage of its accretion. Finally, we consider the role of the forming star in the context of habitability, including effects of the shifting parameters of luminosity, emissions of ionizing radiation, magnetic field interactions, and overall variability, on an evolving Earth-like world. This stellar activity could potentially drive planetary atmospheric loss or drive prebiotic chemistry. In summary, the complex interplay of a diverse array of parameters and processes must be considered as we build our understanding of planetary habitability in the galaxy.

Footnotes

Authors' contributions

All authors contributed to the original draft, the order reflects the extent of the contribution to the draft and editing. A. Segura, R. Smith reviewed and edited the draft. A. Segura supervised all editions, organized the contributions and provided the structure of the draft. All authors reviewed and approved the final version of the manuscript.

Author Disclosure Statement

No competing financial interests exist.

Funding Information

A.S. acknowledges the support of UNAM-PAPIIT IN115524 and IN114026. A.S. and E.S. acknowledge the support of the Virtual Planetary Laboratory, a member of the NASA Nexus for Exoplanet System Science (NExSS), funded via the NASA Astrobiology Program grant No. 80NSSC23K1398. R.S. acknowledges support by the NASA Emerging Worlds program (NNX17AE34G). R.S.’s observations were taken on Maunakea at the Infrared Telescope Facility, which is operated by the University of Hawaii under contract 80HQTR24DA010 with the National Aeronautics and Space Administration. M.T. was supported by NSF CAREER GLOW (AST-2239225) and NASA’s Interdisciplinary Consortia for Astrobiology Research (NNH19ZDA001N-ICAR) under grant number 80NSSC21K0597. E.S. additionally acknowledges support from the Alternative Earths team, funded by NASA Interdisciplinary Consortia for Astrobiology Research (ICAR) program via grant No. 80NSSC21K0594. S.R. thanks UA for support via startup funds.

Associate Editor: Sherry L. Cady

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