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The globular cluster M80. Stars in globular clusters are mainly older metal-poor members of population II.

In astronomy, metallicity is the abundance of elements present in an object that are heavier than hydrogen and helium. Most of the normal currently detectable (i.e. non-dark) matter in the universe is either hydrogen or helium, and astronomers use the word metals as convenient shorthand for all elements except hydrogen and helium. This word-use is distinct from the conventional chemical or physical definition of a metal as an electrically conducting element. Stars and nebulae with relatively high abundances of heavier elements are called metal-rich when discussing metallicity, even though many of those elements are called nonmetals in chemistry.

Metals in early spectroscopy

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Solar spectrum with Fraunhofer lines as it appears visually.

In 1802, William Hyde Wollaston[1] noted the appearance of a number of dark features in the solar spectrum.[2] In 1814, Joseph von Fraunhofer independently rediscovered the lines and began to systematically study and measure their wavelengths, and they are now called Fraunhofer lines. He mapped over 570 lines, designating the most prominent with the letters A through K and weaker lines with other letters.[3][4][5]

About 45 years later, Gustav Kirchhoff and Robert Bunsen[6] noticed that several Fraunhofer lines coincide with characteristic emission lines identified in the spectra of heated chemical elements.[7] They inferred that dark lines in the solar spectrum are caused by absorption by chemical elements in the solar atmosphere.[8] Their observations[9] were in the visible range where the strongest lines come from metals such as sodium, potassium, and iron.[10] In the early work on the chemical composition of the sun the only elements that were detected in spectra were hydrogen and various metals,[11]: 23–24  with the term metallic frequently used when describing them.[11]: Part 2  In contemporary usage in astronomy all the extra elements beyond just hydrogen and helium are termed metallic.

Origin of metallic elements

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See also: Stellar nucleosynthesis and Big Bang nucleosynthesis

The presence of heavier elements results from stellar nucleosynthesis, where the majority of elements heavier than hydrogen and helium in the Universe (metals, hereafter) are formed in the cores of stars as they evolve. Over time, stellar winds and supernovae deposit the metals into the surrounding environment, enriching the interstellar medium and providing recycling materials for the birth of new stars. It follows that older generations of stars, which formed in the metal-poor early Universe, generally have lower metallicities than those of younger generations, which formed in a more metal-rich Universe.

Stellar populations

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Population I star Rigel with reflection nebula IC 2118

Observed changes in the chemical abundances of different types of stars, based on the spectral peculiarities that were later attributed to metallicity, led astronomer Walter Baade in 1944 to propose the existence of two different populations of stars.[12] These became commonly known as population I (metal-rich) and population II (metal-poor) stars. A third, earliest stellar population was hypothesized in 1978, known as population III stars.[13][14][15] These "extremely metal-poor" (XMP) stars are theorized to have been the "first-born" stars created in the Universe.

Common methods of calculation

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Astronomers use several different methods to describe and approximate metal abundances, depending on the available tools and the object of interest. Some methods include determining the fraction of mass that is attributed to gas versus metals, or measuring the ratios of the number of atoms of two different elements as compared to the ratios found in the Sun.

Mass fraction

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Stellar composition is often simply defined by the parameters X, Y, and Z. Here X represents the mass fraction of hydrogen, Y is the mass fraction of helium, and Z is the mass fraction of all the remaining chemical elements. Thus

In most stars, nebulae, H II regions, and other astronomical sources, hydrogen and helium are the two dominant elements. The hydrogen mass fraction is generally expressed as where M is the total mass of the system, and is the mass of the hydrogen it contains. Similarly, the helium mass fraction is denoted as The remainder of the elements are collectively referred to as "metals", and the mass fraction of metals is calculated as

For the surface of the Sun (symbol ), these parameters are measured to have the following values:[16]

Description Solar value
Hydrogen mass fraction
Helium mass fraction
Metal mass fraction

Due to the effects of stellar evolution, neither the initial composition nor the present day bulk composition of the Sun is the same as its present-day surface composition.

Chemical abundance ratios

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The overall stellar metallicity is conventionally defined using the total hydrogen content, since its abundance is considered to be relatively constant in the Universe, or the iron content of the star, which has an abundance that is generally linearly increasing in time in the Universe.[17] Hence, iron can be used as a chronological indicator of nucleosynthesis. Iron is relatively easy to measure with spectral observations in the star's spectrum given the large number of iron lines in the star's spectra[18] (even though oxygen is the most abundant heavy element – see metallicities in H II regions below). The abundance ratio is the common logarithm of the ratio of a star's iron abundance compared to that of the Sun and is calculated thus:[19]

where and are the number of iron and hydrogen atoms per unit of volume respectively, is the standard symbol for the Sun, and for a star (often omitted below). The unit often used for metallicity is the dex, contraction of "decimal exponent".[20] By this formulation, stars with a higher metallicity than the Sun have a positive common logarithm, whereas those more dominated by hydrogen have a corresponding negative value. For example, stars with a value of +1 have 10 times the metallicity of the Sun (10+1); conversely, those with a value of −1 have ⁠1/10⁠, while those with a value of 0 have the same metallicity as the Sun, and so on.[21]

Young population I stars have significantly higher iron-to-hydrogen ratios than older population II stars. Primordial population III stars are estimated to have metallicity less than −6, a millionth of the abundance of iron in the Sun.[22][23] The same notation is used to express variations in abundances between other individual elements as compared to solar proportions. For example, the notation represents the difference in the logarithm of the star's oxygen abundance versus its iron content compared to that of the Sun. In general, a given stellar nucleosynthetic process alters the proportions of only a few elements or isotopes, so a star or gas sample with certain values may well be indicative of an associated, studied nuclear process.

Photometric colors

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Astronomers can estimate metallicities through measured and calibrated systems that correlate photometric measurements and spectroscopic measurements (see also Spectrophotometry). For example, the Johnson UVB filters can be used to detect an ultraviolet (UV) excess in stars,[24] where a smaller UV excess indicates a larger presence of metals that absorb the UV radiation, thereby making the star appear "redder".[25][26][27] The UV excess, δ(U−B), is defined as the difference between a star's U and B band magnitudes, compared to the difference between U and B band magnitudes of metal-rich stars in the Hyades cluster.[28] Unfortunately, δ(U−B) is sensitive to both metallicity and temperature: If two stars are equally metal-rich, but one is cooler than the other, they will likely have different δ(U−B) values[28] (see also Blanketing effect[29][30]). To help mitigate this degeneracy, a star's B−V color index can be used as an indicator for temperature. Furthermore, the UV excess and B−V index can be corrected to relate the δ(U−B) value to iron abundances.[31][32][33]

Other photometric systems that can be used to determine metallicities of certain astrophysical objects include the Strӧmgren system,[34][35] the Geneva system,[36][37] the Washington system,[38][39] and the DDO system.[40][41]

Metallicities in various astrophysical objects

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Stars

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At a given mass and age, a metal-poor star will be slightly warmer. Population II stars' metallicities are roughly ⁠1/1000⁠ to ⁠1/10⁠ of the Sun's but the group appears cooler than population I overall, as heavy population II stars have long since died. Above 40 solar masses, metallicity influences how a star will die: Outside the pair-instability window, lower metallicity stars will collapse directly to a black hole, while higher metallicity stars undergo a type Ib/c supernova and may leave a neutron star.

Relationship between stellar metallicity and planets

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A star's metallicity measurement is one parameter that helps determine whether a star may have a giant planet, as there is a direct correlation between metallicity and the presence of a giant planet. Measurements have demonstrated the connection between a star's metallicity and gas giant planets, like Jupiter and Saturn. The more metals in a star and thus its planetary system and protoplanetary disk, the more likely the system may have gas giant planets. Current models show that the metallicity along with the correct planetary system temperature and distance from the star are key to planet and planetesimal formation. For two stars that have equal age and mass but different metallicity, the less metallic star is bluer. Among stars of the same color, less metallic stars emit more ultraviolet radiation. The Sun, with eight planets and nine consensus dwarf planets, is used as the reference, with a of 0.00.[42][43][44][45][46]

H II regions

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Young, massive and hot stars (typically of spectral types O and B) in H II regions emit UV photons that ionize ground-state hydrogen atoms, knocking electrons free; this process is known as photoionization. The free electrons can strike other atoms nearby, exciting bound metallic electrons into a metastable state, which eventually decay back into a ground state, emitting photons with energies that correspond to forbidden lines. Through these transitions, astronomers have developed several observational methods to estimate metal abundances in H II regions, where the stronger the forbidden lines in spectroscopic observations, the higher the metallicity.[47][48] These methods are dependent on one or more of the following: the variety of asymmetrical densities inside H II regions, the varied temperatures of the embedded stars, and/or the electron density within the ionized region.[49][50][51][52]

Theoretically, to determine the total abundance of a single element in an H II region, all transition lines should be observed and summed. However, this can be observationally difficult due to variation in line strength.[53][54] Some of the most common forbidden lines used to determine metal abundances in H II regions are from oxygen (e.g. [OII] λ = (3727, 7318, 7324) Å, and [OIII] λ = (4363, 4959, 5007) Å), nitrogen (e.g. [NII] λ = (5755, 6548, 6584) Å), and sulfur (e.g. [SII] λ = (6717, 6731) Å and [SIII] λ = (6312, 9069, 9531) Å) in the optical spectrum, and the [OIII] λ = (52, 88) μm and [NIII] λ = 57 μm lines in the infrared spectrum. Oxygen has some of the stronger, more abundant lines in H II regions, making it a main target for metallicity estimates within these objects. To calculate metal abundances in H II regions using oxygen flux measurements, astronomers often use the R23 method, in which

where is the sum of the fluxes from oxygen emission lines measured at the rest frame λ = (3727, 4959 and 5007) Å wavelengths, divided by the flux from the Balmer series Hβ emission line at the rest frame λ = 4861 Å wavelength.[55] This ratio is well defined through models and observational studies,[56][57][58] but caution should be taken, as the ratio is often degenerate, providing both a low and high metallicity solution, which can be broken with additional line measurements.[59] Similarly, other strong forbidden line ratios can be used, e.g. for sulfur, where[60]

Metal abundances within H II regions are typically less than 1%, with the percentage decreasing on average with distance from the Galactic Center.[53][61][62][63][64]

See also

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References

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Further reading

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Revisions and contributorsEdit on WikipediaRead on Wikipedia

Metallicity

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Metallicity in astronomy refers to the abundance of chemical elements heavier than hydrogen and helium within stars, galaxies, interstellar media, and other celestial objects, where these elements—collectively termed "metals" for historical reasons—are produced primarily through stellar nucleosynthesis. This fraction, denoted as the mass fraction Z, represents the proportion of an object's total mass composed of these heavier elements, typically ranging from near zero in the early universe to about 0.014 in the Sun. Metallicity is most commonly quantified using the logarithmic scale [Fe/H], which measures the iron-to-hydrogen ratio relative to solar values, serving as a proxy for overall metal content due to iron's prevalence in stellar spectra. The concept is fundamental to understanding cosmic chemical evolution, as metallicity traces the buildup of heavy elements over time: primordial gas from the was nearly metal-free (Z ≈ 0), but explosions and stellar winds from successive generations of stars enrich the , increasing Z in younger populations. Low-metallicity environments, such as those in ultra-metal-poor stars with [Fe/H] < -4, indicate ancient origins close to the universe's first stars, while higher values correlate with ongoing star formation in mature galaxies. This evolution is evident in the age-metallicity relation observed in stellar populations, where older stars exhibit lower metallicities, providing a chronological map of galactic history. Metallicity influences key astrophysical phenomena, including star formation rates, as higher metal content enhances cooling efficiency in molecular clouds, facilitating the collapse into new stars. In exoplanet systems, stars with elevated metallicity—particularly [Fe/H] > 0—are more likely to host giant planets, due to the availability of solid materials for formation. Additionally, it affects galactic dynamics and observability; metal-rich environments produce more dust, which can obscure light, while metal-poor galaxies, often dwarf or high-redshift systems, emit stronger . Measurements rely on spectroscopic analysis of absorption or emission lines, calibrated against models of stellar atmospheres, with uncertainties typically around 0.1–0.3 dex for precise surveys.

Definition and Fundamentals

Definition of Metallicity

In astronomy, metallicity denotes the proportion of an object's consisting of chemical elements heavier than , which are collectively referred to as "metals" regardless of their in chemistry. This usage contrasts sharply with the chemical definition of metals as elements characterized by properties such as malleability, , and electrical conductivity; in , the term encompasses non-metals like carbon, oxygen, and alongside true metals. Elements with s greater than 2 are included, as ( 1) and ( 2) dominate the composition of most astronomical bodies, with the remainder tracing past nucleosynthetic processes. The standard measure of metallicity is the mass fraction , defined as the ratio of the mass of all metals to the total of the object: Z=MmetalsMtotalZ = \frac{M_{\text{metals}}}{M_{\text{total}}}. Complementary mass fractions are used for (X) and (Y), satisfying the relation X+Y+Z=1X + Y + Z = 1. For the Sun, recent 3D non-local spectroscopic analyses yield a photospheric metallicity of Z⊙≈0.0139±0.0006Z_\odot \approx 0.0139 \pm 0.0006 (as of 2021), though alternative determinations considering helioseismic constraints suggest values up to approximately 0.017; this is lower than older estimates around 0.019–0.02. Because directly measuring the total metal content is challenging, astronomers often employ proxies such as the abundance of iron-peak elements (e.g., iron, , ), which form in similar nucleosynthetic environments and correlate strongly with overall metallicity. The helium abundance Y is particularly complementary, as it reflects primordial contributions from augmented by stellar processing, while Z captures enrichment from heavier elements. The term "metallicity" entered astronomical usage in the early , particularly during the , as revealed composition patterns in stars, even though many designated "metals" lack metallic properties.

Importance in Astrophysics

Metallicity plays a crucial role in stellar physics by influencing opacity, which governs transport within . Higher metallicity increases opacity due to the presence of more electrons from ionized metals, leading to slower convective and radiative transport, cooler surface temperatures, and extended main-sequence lifetimes for a given . Conversely, low-metallicity exhibit reduced opacity, enabling more efficient escape, resulting in hotter effective temperatures and bluer colors on the Hertzsprung-Russell diagram; for a fixed , these must be more massive to compensate for their higher temperatures and faster nuclear burning rates. This metallicity dependence also affects rotational evolution and loss through stellar winds, with low-metallicity massive experiencing weaker winds and retaining more . In the context of chemical evolution, metallicity serves as a tracer of the universe's enrichment , reflecting the cumulative impact of over cosmic time. As stars form from interstellar gas, they synthesize and eject heavy elements via supernovae and stellar winds, progressively increasing the metallicity of subsequent generations; thus, higher metallicity correlates with later epochs and more advanced stages of galactic chemical evolution. This process links efficiency, inflow and outflow rates, and the , providing insights into how galaxies transition from metal-poor, primordial conditions to metal-enriched states observed today. Metallicity provides essential observational diagnostics for interpreting galactic spectra and modeling formation processes. It modulates the strength and profiles of emission and absorption lines—such as those from oxygen, , and iron—allowing astronomers to infer states, dust content, and excitation conditions in the ; these line ratios are calibrated into diagnostics like the R23 index for estimating gas-phase abundances. In galaxy formation models, metallicity gradients and dispersions help constrain feedback mechanisms, merger histories, and rates, enabling age-dating of stellar populations through integrated light analysis. On cosmological scales, the mean metallicity of the —defined as the total mass in metals divided by total baryonic mass—has evolved from values below 10^{-3} Z_⊙ at z ≈ 10 to near-solar levels today, driven by metal enrichment from successive generations of stars and galaxies. Recent observations (as of 2025) reveal that individual galaxies at z ≈ 10 can exhibit metallicities ranging from ∼0.01 to 0.1 Z_⊙, indicating rapid early enrichment, though the cosmic average remains low due to vast unenriched intergalactic medium. This evolution interconnects with by influencing the cooling efficiency of primordial gas, facilitating earlier , and with feedback processes, where metal-rich outflows regulate subsequent accretion and prevent over-enrichment. For instance, in unresolved galaxies at high , integrated metallicity serves as a proxy for history, as higher abundances indicate prolonged or intense episodes of nucleosynthetic enrichment relative to gas dilution.

Historical Context

Metals in Early Spectroscopy

The discovery of absorption lines in the solar spectrum marked a pivotal moment in early stellar . In 1814, systematically observed and cataloged hundreds of dark lines crossing the continuous spectrum of the Sun, which became known as and served as foundational standards for spectral analysis. These observations revealed intricate patterns but initially lacked explanation regarding their origins. Building on this, in 1859, and conducted laboratory experiments with flames and spectroscopes, identifying key —such as the sodium D lines and iron absorption features—as arising from the same chemical elements present in the Sun's atmosphere. Their work demonstrated that the Sun contained terrestrial substances like sodium and iron, fundamentally linking laboratory chemistry to celestial phenomena and establishing as a tool for remote elemental detection. A notable early highlight came in 1868 when , during observations of a , detected an unknown yellow line at 587.49 nm in the , which did not match any known terrestrial element. Lockyer interpreted this as evidence of a new heavy element, which he named (from the Greek helios for sun), initially classifying it as a metal due to its spectral behavior and the prevailing convention of naming metallic elements with the suffix "-ium." This discovery, independently corroborated by Pierre Janssen's eclipse observations, expanded the periodic table beyond Earth-bound samples and underscored the potential for stellar spectra to reveal novel cosmic chemistry, though 's true nature as a was only confirmed decades later on in 1895. In the late 19th and early 20th centuries, astronomers adopted the term "metals" to describe the prominent bright or absorption lines from elements such as calcium (e.g., the H and K lines), sodium, and iron dominating stellar spectra, reflecting their chemical classification as metals on . Early analyses assumed that cosmic abundances mirrored terrestrial compositions, with iron and other heavy elements expected to be prevalent based on geological samples and meteorites. Henry Norris Russell advanced this in the 1920s by compiling comprehensive tables of elemental abundances derived from solar and stellar spectra, positing that most stars shared a solar-like composition rich in these "metals." These tables, published in works like his 1929 paper on the Sun's atmosphere, provided quantitative estimates but relied on curve-of-growth methods that often overlooked ionization effects. Such assumptions faced initial challenges from emerging evidence of abundance variations. In her 1925 doctoral thesis, Stellar Atmospheres, Cecilia Payne-Gaposchkin applied Meghnad Saha's theory to analyze strengths across stellar types, concluding that and dominated stellar compositions far more than in terrestrial rocks, with "metals" like iron being far less abundant than previously thought. This work highlighted discrepancies between astrophysical and Earth-based abundances, attributing them to differences in and states rather than uniform cosmic chemistry, though her findings were initially met with skepticism. Payne-Gaposchkin's analysis thus offered the first rigorous hints that stellar "metallicity"—the relative proportion of elements heavier than —could vary, laying groundwork for refined concepts in later decades.

Evolution of Metallicity Concepts

Following the initial discoveries of metallic lines in stellar spectra in the early , the concept of metallicity evolved significantly after 1925, transitioning from absolute abundance estimates to relative measurements normalized against the Sun. In the 1940s, Albrecht Unsöld pioneered the adoption of logarithmic scales, defining [Fe/H] as log₁₀ (N_Fe / N_H)star - log₁₀ (N_Fe / N_H)⊙, which facilitated comparisons across by emphasizing deviations from solar values rather than absolute numbers. This shift addressed uncertainties in atomic data and model atmospheres, making metallicity a practical proxy for overall heavy-element content. Key milestones in the mid-20th century further refined these concepts. Lawrence H. Aller and Dean B. McLaughlin contributed to solar abundance revisions in the 1940s through detailed spectroscopic analyses, updating values for iron and other metals based on improved line identifications and oscillator strengths. By the , observations revealed enhancements in alpha-elements (such as Mg, Si, Ca, and ) relative to iron in old, metal-poor stars, indicating distinct nucleosynthetic contributions from massive stars and Type II supernovae early in galactic history. These findings, initially noted in halo population studies, underscored metallicity as a chronological tracer of stellar generations. In the , solar metallicity was commonly taken as Z_⊙ ≈ 0.02, based on contemporaneous photospheric models. Subsequent updates lowered this value amid refining techniques and data. The 1989 compilation by Edward Anders and Martin Grevesse set Z_⊙ ≈ 0.019, incorporating meteoritic constraints. Further revisions by Martin Asplund et al. in 2009 and 2021, using 3D hydrodynamical models and non-LTE corrections, reduced Z_⊙ to approximately 0.014 (specifically 0.0134 in 2009 and 0.0139 ± 0.0006 in 2021), driven by discrepancies between low-metallicity solar models and helioseismic observations of sound speeds and abundance in the solar interior. These changes highlighted tensions in opacity calculations and modeling. In the 2020s, ongoing refinements from advanced 3D non-local thermodynamic equilibrium simulations continue to adjust individual element abundances, with Z_⊙ stabilizing around 0.014 while resolving some spectroscopic uncertainties. Complementing this, (JWST) observations of primitive, low-metallicity systems in the early (z > 10) provide empirical benchmarks for low-[Fe/H] regimes, informing calibrations of metallicity scales and revealing alpha-enhancements in metal-poor environments akin to those in ancient stars. These data bridge conceptual gaps in the low-metallicity tail, enhancing the universality of relative abundance frameworks.

Origin of Metals

Big Bang Nucleosynthesis Limitations

(BBN) occurred in the first few minutes after the , when the cooled to temperatures around 0.1 MeV, allowing protons and neutrons to combine into light nuclei. This process primarily synthesized (about 75% by mass), (about 25% by mass), and trace amounts of , , lithium-7, and beryllium-7, with no significant production of heavier elements. The neutron-to-proton ratio froze out at approximately 1:6 during the rapid expansion at around 1 MeV, limiting the available neutrons for fusion beyond helium-4. Most neutrons were incorporated into 4^4He nuclei via the reaction p+n→2p + n \to ^2H + \gammafollowedbyfollowed by^2H+2H + ^2H \to ^4He+2γHe + 2\gamma, leaving insufficient free neutrons for further captures. Subsequent neutron capture processes stalled due to the instability of intermediate nuclei like 8^8Be and the low density of the expanding universe, which prevented the buildup of elements beyond lithium. Observations confirm these predictions, with the primordial helium mass fraction Yp≈0.24−0.25Y_p \approx 0.24-0.25 derived from cosmic microwave background (CMB) measurements of the baryon density and corroborated by deuterium surveys. The initial metallicity Zinitial≈0Z_{\rm initial} \approx 0, as BBN produced no metals (elements heavier than helium), with all subsequent enrichment occurring post-BBN through stellar processes. BBN also predicts a primordial deuterium-to-hydrogen ratio D/H∼10−5D/H \sim 10^{-5}, matching observations in metal-poor quasar absorption systems that probe nearly pristine gas. This primordial composition, with zero metallicity, set the conditions for the formation of the first stars (Population III), which assembled in dark matter halos at redshifts z≳20z \gtrsim 20. Without metals to facilitate fine-structure line cooling, the gas relied on molecular hydrogen for cooling, resulting in higher collapse temperatures and Jeans masses that favored the formation of massive stars (10−1000M⊙10-1000 M_\odot). These stars initiated the chemical enrichment of the universe.

Stellar and Supernova Nucleosynthesis

Stellar nucleosynthesis begins with hydrogen and helium fusion in the cores of stars, producing heavier elements through successive stages of nuclear burning. In massive stars (typically greater than 8 solar masses), the CNO cycle dominates the production of carbon, nitrogen, and oxygen by catalyzing hydrogen fusion via proton captures on these seed nuclei, recycling them in a cycle that releases energy while enhancing CNO abundances in the stellar envelopes. Asymptotic giant branch (AGB) stars, with initial masses of 1–8 solar masses, contribute to the slow neutron-capture process (s-process), where neutrons are captured on iron-peak seeds at rates slow enough for beta decays to occur between captures, synthesizing elements like barium (Ba) and strontium (Sr) in their convective thermal pulse phases. The rapid neutron-capture process (r-process), responsible for heavy elements such as europium (Eu) and gold (Au), occurs in extreme neutron-rich environments, primarily from neutron star mergers where dynamically ejected material undergoes intense neutron bombardment, though core-collapse supernovae may contribute in some models. Supernovae play a crucial role in dispersing these synthesized metals into the , fueling galactic chemical evolution. Core-collapse supernovae (Type II) from massive stars (>8 solar masses) explode upon core iron formation and collapse, ejecting alpha-elements like oxygen (O), magnesium (Mg), and silicon (Si) produced in pre-explosion silicon and oxygen burning shells, with yields enriched in these species relative to iron. Type Ia supernovae, arising from in binary systems reaching the , contribute primarily to iron-peak elements (e.g., Fe, Ni) through explosive carbon-oxygen burning in a thermonuclear runaway, adding significant iron enrichment on longer timescales due to their longer progenitor lifetimes. The cumulative effect of these processes on interstellar metallicity is described by chemical models, where the final metallicity ZfinalZ_{\rm final} integrates the metal yield y(Z)y(Z) from stellar populations—dependent on the initial metallicity—with the rate ψ(t)\psi(t), divided by the gas surface Σgas\Sigma_{\rm gas}: Zfinal=∫y(Z)ψ(t) dtΣgas.Z_{\rm final} = \frac{\int y(Z) \psi(t) \, dt}{\Sigma_{\rm gas}}. This instantaneous approximation assumes prompt return of metals to the gas phase, with typical yields of several (~2–6) es of metals per core-collapse event, scaling with the to an effective yield of approximately 0.02–0.03 es of metals per of stars formed. The earliest metals originated from Population III stars, the metal-free first generation formed at redshifts z≈20–30z \approx 20–30, whose deaths as pair-instability supernovae in the mass range 140–260 solar masses ejected 100–300 solar masses of metals per event, primarily oxygen and from explosive helium and carbon burning, without leaving remnants. Recent observations from 2022–2025 have detected metal enrichment signatures, such as oxygen emission lines, in galaxies at z>10z > 10, confirming rapid early chemical evolution consistent with these primordial explosions. While produced negligible metals beyond , these stellar processes initiated the transition to metal-enriched subsequent generations.

Measurement Techniques

Mass Fraction and Abundance Ratios

Metallicity is quantified as the mass fraction ZZ, defined as the total mass of elements heavier than (collectively termed "metals" in ) divided by the total of the composition, Z=∑i>2mi/MtotalZ = \sum_{i>2} m_i / M_{\rm total}, where mim_i are the masses of individual metal and MtotalM_{\rm total} is the overall . This fraction arises primarily from spectroscopic analysis of stellar or gaseous atmospheres, where line strengths are interpreted using models of atomic processes to derive contributions. The solar metallicity is Z⊙=0.0139±0.0006Z_\odot = 0.0139 \pm 0.0006, reflecting the photospheric composition derived from advanced 3D hydrodynamical simulations and non-local (non-LTE) . However, recent helioseismic analyses as of 2024 suggest potentially higher solar metallicities (Z/X ≈ 0.0225) to better match internal structure constraints. The mass fraction ZZ relates directly to the (XX) and (YY) mass fractions via the identity Z=1−X−YZ = 1 - X - Y, assuming negligible contributions from other primordial elements. For the Sun, typical values are X⊙≈0.744X_\odot \approx 0.744 and Y⊙≈0.242Y_\odot \approx 0.242, yielding the quoted Z⊙Z_\odot. In deriving these fractions from observations, the plays a central role by balancing the populations of ionized stages in , nr+1nenr=2gr+1gr(2πmekTh2)3/2e−χr/kT\frac{n_{r+1} n_e}{n_r} = \frac{2 g_{r+1}}{g_r} \left( \frac{2\pi m_e k T}{h^2} \right)^{3/2} e^{-\chi_r / kT}, where nn denotes number densities, gg statistical weights, χr\chi_r ionization potentials, and other terms are standard constants; this allows conversion of observed equivalent widths to elemental number densities, and subsequently to mass fractions assuming known atomic masses. Abundance ratios provide a normalized measure of relative content, expressed logarithmically as [X/Y]=log⁡10((NX/NY)obj(NX/NY)⊙)[\mathrm{X/Y}] = \log_{10} \left( \frac{(N_{\mathrm{X}}/N_{\mathrm{Y}})_{\rm obj}}{(N_{\mathrm{X}}/N_{\mathrm{Y}})_\odot} \right), where NN is the of X and Y in the object and Sun, respectively. The iron-to-hydrogen ratio [Fe/H][\mathrm{Fe/H}] serves as a common proxy for overall metallicity due to the prevalence and strength of iron lines across stellar spectra, capturing the enrichment history while approximating ZZ to within typical uncertainties. Enhancements in alpha elements (O, Ne, Mg, Si, S, Ca, Ti) relative to iron, denoted [α/Fe][\alpha/\mathrm{Fe}], indicate contributions from core-collapse supernovae in early galactic evolution, often elevated by 0.3–0.4 dex in metal-poor populations. Uncertainties in derived abundances stem partly from atomic data, with errors in oscillator strengths (log⁡gf\log gf) contributing approximately 0.1–0.2 dex to individual elemental ratios. Modern refinements, including non-LTE line formation and 3D atmospheric granulation effects, adjust the solar [Fe/H] zero point upward by about 0.02 dex relative to classical 1D local thermodynamic equilibrium (LTE) models.

Photometric and Spectroscopic Methods

Spectroscopic methods provide one of the most direct ways to measure metallicity by analyzing absorption lines in stellar spectra, particularly those of iron, which serve as proxies for overall metal content. Line-strength indices, such as the equivalent width (EW) of iron (Fe) lines divided by the local continuum level, quantify the depth of these lines relative to the surrounding spectrum, allowing estimation of iron abundance [Fe/H] with typical precisions of 0.1–0.2 dex for moderate-resolution spectra (R ≈ 20,000). For higher precision, high-resolution spectroscopy (R > 20,000) enables detailed curve-of-growth analysis, where the EW of multiple lines is plotted against their expected oscillator strengths to derive abundances by fitting theoretical curves that account for damping, thermal broadening, and saturation effects, achieving [Fe/H] uncertainties around 0.1 dex. This technique has been widely applied in surveys like the DESI Early Data Release, where data-driven methods using differential spectra yield abundances for thousands of stars with minimal systematic biases. Photometric methods infer metallicity indirectly through correlations between broadband colors or specific filter passbands and metal abundance, offering efficiency for large samples despite lower precision (typically 0.2–0.5 dex). In color-magnitude diagrams, metal-rich stars appear redder due to increased line blanketing in the blue, with indices like (B–V) correlating with [Fe/H] via calibrated relations derived from spectroscopic benchmarks. Metallicity-sensitive filters, such as the Ca II H and K (CaHK) index, target ultraviolet absorption lines of calcium, which strengthen with higher metallicity; this narrowband photometry, often combined with (b–y) colors in Strömgren systems, provides [Fe/H] estimates for evolved stars with scatter reduced to ~0.15 dex when calibrated against high-resolution data. Beyond traditional and photometry, complementary techniques leverage indirect probes of stellar interiors or all-sky surveys for metallicity determination. Asteroseismology analyzes oscillations in stars like red giants, where the large separation δν scales approximately as δν ∝ (1/√μ)^{1/2} with mean molecular weight μ influenced by metallicity through its effect on speed and mean , enabling [Fe/H] constraints to ~0.2 dex when combined with . The Data Release 3 (2022) extends photometric metallicities to approximately 220 million stars using low-resolution XP spectra convolved with synthetic filters, deriving [Fe/H] for ~694,000 giants with median precision of 0.13 dex via calibrations against spectroscopic catalogs. Additionally, the APOGEE survey (2010s–2020s) employs near-infrared H-band (R ≈ 22,500) to penetrate dust-obscured regions, measuring [Fe/H] for hundreds of thousands of stars in the bulge with typical uncertainties of 0.1 dex by fitting synthetic spectra to lines of multiple elements.

Metallicity in Stars

Stellar Metallicity Determination

Stellar metallicity is primarily determined through tailored to the stellar type, with high-dispersion (resolution R ≥ 25,000) commonly applied to FGK dwarfs to measure iron abundances via equivalent widths of lines or synthetic fitting using model atmospheres like MARCS. This method achieves precisions of 0.04–0.06 dex in high signal-to-noise (S/N ≥ 50) observations, as seen in surveys employing tools such as FERRE or SME, though non-local (non-LTE) corrections for iron lines are essential to mitigate biases at low metallicities ([Fe/H] < -0.5 dex). For red giants, low-resolution (R ~ 1,000–5,000) in the near-infrared K-band is often preferred due to reduced line crowding and sensitivity to molecular features, enabling metallicity estimates from empirical calibrations against effective temperatures and surface gravities. These determinations are calibrated using open clusters with well-established solar-like compositions, such as M67, which has an average [Fe/H] = 0.023 ± 0.015 derived from high-resolution spectra of main-sequence and giant members. However, challenges arise in cool stars (T_eff < 5,000 K), where severe line blending in crowded spectra leads to overestimation of equivalent widths and thus inflated metallicities, particularly for M dwarfs and giants. Additionally, granulation effects—convective surface motions—introduce variability in line profiles that correlates with metallicity, with metal-poor stars exhibiting reversed granulation (cooler intergranular regions) that can bias abundance measurements by up to 0.1–0.2 dex without 3D atmospheric modeling. In the Galactic disk, radial metallicity gradients further complicate interpretations, with d[Fe/H]/dR ≈ -0.042 ± 0.011 dex/kpc observed for thin-disk populations at heights 1.3–1.7 kpc above the plane, reflecting inside-out formation and migration effects. Large-scale surveys have revolutionized stellar metallicity mapping by providing homogeneous [Fe/H] data for millions of stars. The LAMOST survey has delivered low-resolution spectra for over 10 million stars, yielding [Fe/H] estimates with typical precisions of 0.25–0.3 dex through data-driven pipelines like LSP3, enabling studies of disk structure despite limitations in cool-star accuracy. The GALAH DR3 (2021) extends this with medium-resolution (R ~ 28,000) observations of nearly 600,000 stars, reporting [X/Fe] ratios for 30 elements (including non-LTE corrections for 11) across nucleosynthetic pathways, achieving 0.05–0.1 dex precision for iron and key alphas. More recently, the DESI survey's DR1 (2024–2025) includes a stellar catalog from low-resolution spectra of over 10 million Milky Way targets, providing [Fe/H] and radial velocities to map halo substructures and gradients with ~0.2 dex accuracy. Extreme cases, such as hyper metal-poor stars with [Fe/H] < -4, test the limits of these methods and probe early enrichment. The star SMSS J0313-6708, discovered in 2014 via the SkyMapper Southern Survey, holds the record with [Fe/H] ≈ -7.8 (upper limit from non-detections in high-resolution spectra), its light-element enhancements (C, N, Mg) attributed to a single low-energy supernova from a Population III progenitor.

Relation to Planet Formation

In the core-accretion model of planet formation, higher stellar metallicity, parameterized as [Fe/H], enhances the availability of solid materials in the protoplanetary disk, facilitating the growth of rocky planetary cores. This abundance of refractory elements and s allows cores to reach the critical mass necessary for runaway gas accretion, leading to formation. Theoretical models predict a metallicity threshold of [Fe/H] > −0.5 below which formation is inefficient, as lower metallicity disks struggle to assemble sufficiently massive cores within the disk lifetime. The , the radial boundary beyond which water ice condenses, further influences solid inventory by increasing the surface density of buildable materials interior to it, amplifying metallicity effects on core growth. Observational surveys using (RV) and transit methods, such as those from the past two decades, reveal distinct metallicity dependencies for different types. Super-Earths and sub-Neptunes form across a broad range of host star metallicities, including subsolar values down to [Fe/H] ≈ −0.5, indicating that these smaller planets require less solid material and can assemble efficiently even in metal-poor disks. In contrast, Jupiter-mass giants show a strong positive with supersolar metallicity ([Fe/H] > 0), with occurrence rates increasing by factors of 3–4 from [Fe/H] = 0 to +0.5, consistent with the core-accretion paradigm. Hot Jupiters exhibit a pronounced "metallicity desert" at low [Fe/H] (< 0), where their detection frequency drops sharply, underscoring the role of enhanced disk solids in enabling inward migration and retention of close-in giants. Studies of solar twins, stars closely matching the Sun's parameters including [Fe/H] ≈ 0, provide insights into how subtle metallicity variations link to circumstellar dust and planet formation. For instance, the solar twin HIP 56948, with [Fe/H] = +0.02, displays a refractory element pattern slightly depleted relative to volatiles compared to the Sun, interpreted as evidence of dust processing or terrestrial planet formation removing solids from the disk. Broader samples of solar twins show correlations between higher refractory abundances and indicators of dust-rich debris disks, suggesting that modest increases in [Fe/H] promote the formation and retention of planetesimals that evolve into planets or observable dust. Recent (JWST) observations from 2023 onward have begun probing exoplanet atmospheres directly, revealing metallicities that align with their host stars' values. Transmission spectroscopy of hot Jupiters like WASP-39b shows atmospheric [Fe/H] consistent with the solar-metallicity host, supporting models where planet bulk composition inherits the disk's metallicity. This matching extends to other targets, with envelope metallicities scaling as Mplanet∝Zdisk×ηM_{\rm planet} \propto Z_{\rm disk} \times \eta, where ZdiskZ_{\rm disk} is the disk metallicity (approximated by stellar [Fe/H]) and η\eta is the accretion efficiency, modulated by factors like the snow line position.

Metallicity in Other Astrophysical Objects

H II Regions and Interstellar Medium

H II regions, which are ionized nebulae surrounding young, massive stars, serve as key laboratories for measuring gas-phase metallicity in the interstellar medium (ISM). The direct method for determining oxygen abundance in these regions relies on electron temperature (T_e) derived from the ratio of auroral to nebular forbidden lines, specifically [O III] λ4363 to [O III] λλ4959,5007, which provides T_e[O III] and enables calculation of the O^{++}/H^+ ionic ratio; total O/H is then obtained by adding O^+/H^+ from [O II] λλ3726,3729 lines, assuming ionization corrections based on models. This temperature-based approach minimizes assumptions about ionization structure and is preferred for its accuracy in low-metallicity environments, yielding 12 + log(O/H) values with uncertainties typically around 0.05–0.1 dex when the weak [O III] λ4363 line is detectable. For regions where the direct T_e method is infeasible due to faint auroral lines, the R_{23} calibration is widely used, defined as log R_{23} = log[ ([O II] λλ3726,3729 + [O III] λλ4959,5007) / Hβ ], which correlates with 12 + log(O/H) in a double-valued manner (lower branch for Z < 0.5 Z_⊙, upper for higher Z), often calibrated against direct measurements and requiring an additional parameter like [O III]/[O II] to resolve the turnover. In the Milky Way, typical H II region metallicities from these methods average 12 + log(O/H) ≈ 8.7, reflecting near-solar abundances in the solar neighborhood. A representative example is the , where direct measurements yield [O/H] ≈ solar (12 + log(O/H) ≈ 8.6), consistent with its proximity to the Sun and recent enrichment by massive stars. In the broader ISM, metallicity is probed using a combination of recombination lines like Hα (for hydrogen density) and forbidden lines such as [N II] λ6584 and [S II] λλ6716,6731 for ionic abundances, with total metallicities derived via photoionization models that account for varying ionization parameters. Dust depletion significantly affects observed abundances, as refractory elements like iron and silicon are incorporated into grains, reducing gas-phase measurements by 0.2–0.5 dex compared to total (gas + dust) values, with oxygen less affected but still showing mild depletion in dense regions. Bayesian photoionization fitting methods, developed in the 2010s, have enabled spatially resolved ISM metallicity mapping by inferring Z and ionization parameter from emission-line ratios, revealing gradients and inhomogeneities in nearby galaxies with precision down to 0.1–0.2 dex. These techniques extend to high-redshift studies, where strong-line methods using far-infrared lines observed by ALMA, such as [C II] λ158 μm relative to CO or far-IR continuum, indicate metallicities around 0.1 Z_⊙ in low-mass galaxies at z ≈ 2, highlighting rapid enrichment in the early universe. In dwarf galaxies, H II region metallicities often exhibit bimodality, with a low-metallicity peak (12 + log(O/H) < 7.5) linked to pristine gas inflows and a higher peak approaching 8.0 from star formation feedback, as mapped in samples like those from .

Galaxies and Stellar Populations

In galaxies, stellar populations are classified based on their age, location, and chemical composition, with metallicity serving as a key discriminator. Population I stars, which are young and form primarily in the disks of spiral galaxies like the , exhibit near-solar iron abundances of approximately [Fe/H] ≈ 0, reflecting ongoing enrichment from multiple generations of stars. In contrast, Population II stars are ancient, metal-poor relics residing in galactic halos and bulges, with typical iron abundances [Fe/H] < -1, indicative of formation in the early universe before significant metal enrichment occurred. These populations arise from distinct nucleosynthetic processes, where Population II stars show α-element enhancement ([α/Fe] ≈ +0.3) due to rapid enrichment by core-collapse supernovae (Type II) that produce α-elements like oxygen and magnesium before Type Ia supernovae contribute iron-peak elements. Metallicity distributions across galaxies often exhibit spatial gradients, providing insights into chemical evolution and dynamical processes. In spiral galaxies, radial metallicity gradients are typically negative, with oxygen abundance decreasing outward at a rate of d[O/H]/dR ≈ -0.04 dex/kpc, resulting from inside-out formation where the inner regions accumulate metals more efficiently through star formation and inflows. Elliptical galaxies, however, can display inverted (positive) gradients in some cases, particularly in low-mass systems, where outer regions may appear more metal-rich due to mergers or differential enrichment. Additionally, the mass-metallicity relation links stellar mass to average metallicity, with dwarf galaxies showing a proportionality M_* ∝ Z, as lower-mass systems retain fewer metals due to stronger outflows and less efficient recycling. Observational studies reveal diverse metallicity patterns in specific systems and across cosmic time. The Milky Way's halo has an average iron abundance of [Fe/H] ≈ -1.7 (as of 2023), with the outer halo peaking at around [Fe/H] ≈ -2.2, dominated by ancient Population II stars accreted from disrupted satellites. Nearby irregular galaxies like the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) have overall metallicities around 0.5 Z_⊙ and 0.2 Z_⊙, respectively, lower than solar due to their lower masses and interaction histories. At high redshifts (z > 6), James Webb Space Telescope (JWST) observations using the NIRSpec instrument have detected galaxies with metallicities ranging from 0.01 to 0.1 Z_⊙, such as a z = 7.20 system with 12 + log(O/H) ≈ 7.64, highlighting rapid early enrichment in the era. A notable puzzle in galactic chemical evolution is the G-dwarf problem, which describes the observed scarcity of low-metallicity ([Fe/H] < -1) G-type dwarf in the solar neighborhood compared to predictions from simple closed-box models of chemical . This discrepancy is largely resolved by radial migration models, where stars from inner, metal-rich regions churn outward over billions of years due to transient spiral arms and bars, flattening the local metallicity distribution function.

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