Astronomy:Metallicity

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Short description: Relative abundance of heavy elements in a star or other astronomical object
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 a convenient short term 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 solid. Stars and nebulae with relatively high abundances of heavier elements are called "metal-rich" in astrophysical terms, even though many of those elements are nonmetals in chemistry.

Origin

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

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.[1] 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.[2][3][4] These "extremely metal-poor" (XMP) stars are theorized to have been the "first-born" stars created in the Universe.

Common methods of calculation

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

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

[math]\displaystyle{ X + Y + Z = 1 }[/math]

In most stars, nebulae, HII regions, and other astronomical sources, hydrogen and helium are the two dominant elements. The hydrogen mass fraction is generally expressed as [math]\displaystyle{ \ X \equiv \tfrac{m_\mathsf{H}}{M}\ , }[/math] where M is the total mass of the system, and [math]\displaystyle{ \ m_\mathsf{H}\ }[/math] is the mass of the hydrogen it contains. Similarly, the helium mass fraction is denoted as [math]\displaystyle{ \ Y \equiv \tfrac{m_\mathsf{He}}{M} ~. }[/math] The remainder of the elements are collectively referred to as "metals", and the metallicity – the mass fraction of elements heavier than helium – is calculated as

[math]\displaystyle{ Z = \sum_{e \gt \mathsf{He}} \tfrac{m_e}{M} = 1 - X - Y ~. }[/math]

For the surface of the Sun (symbol [math]\displaystyle{ \odot }[/math]), these parameters are measured to have the following values:[5]

Description Solar value
Hydrogen mass fraction [math]\displaystyle{ \ X_\odot = 0.7381\ }[/math]
Helium mass fraction [math]\displaystyle{ \ Y_\odot = 0.2485\ }[/math]
Metallicity [math]\displaystyle{ \ Z_\odot = 0.0134\ }[/math]

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

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.[6] 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 (even though oxygen is the most abundant heavy element – see metallicities in HII 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:[7]

[math]\displaystyle{ \left[ \frac{ \mathsf{Fe} }{ \mathsf{H} } \right] ~=~ \log_{10}{\left( \frac{N_{\mathsf{Fe}}}{N_{\mathsf{H}} } \right)_\star } -~ \log_{10}{\left(\frac{N_{ \mathsf{Fe}} }{ N_{\mathsf{H}} } \right)_\odot}\ , }[/math]

where [math]\displaystyle{ \ N_{\mathsf{Fe}}\ }[/math] and [math]\displaystyle{ \ N_{\mathsf{H}}\ }[/math] are the number of iron and hydrogen atoms per unit of volume respectively, [math]\displaystyle{ \odot }[/math] is the standard symbol for the Sun, and [math]\displaystyle{ \star }[/math] for a star (often omitted below). The unit often used for metallicity is the dex, contraction of "decimal exponent". 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 [math]\displaystyle{ \ \bigl[\tfrac{ \mathsf{Fe} }{ \mathsf{H} } \bigr]_\star\ }[/math] value of +1 have 10 times the metallicity of the Sun (10+1); conversely, those with a [math]\displaystyle{ \ \bigl[\tfrac{ \mathsf{Fe} }{ \mathsf{H} } \bigr]_\star\ }[/math] value of −1 have 1/10, while those with a [math]\displaystyle{ \ \bigl[\tfrac{ \mathsf{Fe} }{ \mathsf{H} } \bigr]_\star\ }[/math] value of 0 have the same metallicity as the Sun, and so on.[8]

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.[9][10] The same notation is used to express variations in abundances between other individual elements as compared to solar proportions. For example, the notation [math]\displaystyle{ \ \bigl[\tfrac{ \mathsf{O} }{ \mathsf{Fe} } \bigr]\ }[/math] 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 [math]\displaystyle{ \ \bigl[\tfrac{ \mathsf{?} }{ \mathsf{Fe} } \bigr]_\star\ }[/math] values may well be indicative of an associated, studied nuclear process.

Photometric colors

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,[11] where a smaller UV excess indicates a larger presence of metals that absorb the UV radiation, thereby making the star appear "redder".[12][13][14] 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.[15] 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[15] (see also Blanketing effect[16][17]). 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.[18][19][20]

Other photometric systems that can be used to determine metallicities of certain astrophysical objects include the Strӧmgren system,[21][22] the Geneva system,[23][24] the Washington system,[25][26] and the DDO system.[27][28]

Metallicities in various astrophysical objects

Stars

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 [math]\displaystyle{ \left(\ \bigl[ \tfrac{ \mathsf{Fe} }{ \mathsf{H} } \bigr]\ = {-3.0}\ ...\ {-1.0}\ \right)\ , }[/math] 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

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 [math]\displaystyle{ \ \bigl[\tfrac{ \mathsf{Fe} }{ \mathsf{H} } \bigr]\ }[/math] of 0.00.[29][30][31][32][33]

HII regions

Young, massive and hot stars (typically of spectral types O and B) in HII regions emit UV photons that ionize ground-state hydrogen atoms, knocking electrons and protons 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 HII regions, where the stronger the forbidden lines in spectroscopic observations, the higher the metallicity.[34][35] These methods are dependent on one or more of the following: the variety of asymmetrical densities inside HII regions, the varied temperatures of the embedded stars, and/or the electron density within the ionized region.[36][37][38][39]

Theoretically, to determine the total abundance of a single element in an HII region, all transition lines should be observed and summed. However, this can be observationally difficult due to variation in line strength.[40][41] Some of the most common forbidden lines used to determine metal abundances in HII 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 HII regions, making it a main target for metallicity estimates within these objects. To calculate metal abundances in HII regions using oxygen flux measurements, astronomers often use the R23 method, in which

[math]\displaystyle{ R_{23} = \frac{\ \left[\ \mathsf{O}^\mathsf{II} \right]_{3727~\AA} + \left[\ \mathsf{O}^\mathsf{III} \right]_{4959~\AA + 5007~\AA}\ }{\Bigl[\ \mathsf{ H}_\mathsf{\beta} \Bigr]_{4861 ~\AA} }\ , }[/math]

where [math]\displaystyle{ \ \left[\ \mathsf{O}^\mathsf{II} \right]_{3727~\AA} + \left[\ \mathsf{O}^\mathsf{III} \right]_{4959~\AA + 5007~\AA}\ }[/math] 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.[42] This ratio is well defined through models and observational studies,[43][44][45] 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.[46] Similarly, other strong forbidden line ratios can be used, e.g. for sulfur, where[47]

[math]\displaystyle{ S_{23} = \frac{\ \left[\ \mathsf{S}^\mathsf{II} \right]_{6716~\AA + 6731~\AA} + \left[\ \mathsf{S}^\mathsf{III} \right]_{9069~\AA + 9532~\AA}\ }{\Bigl[\ \mathsf{H}_\mathsf{\beta} \Bigr]_{4861 ~\AA} } ~. }[/math]

Metal abundances within HII regions are typically less than 1%, with the percentage decreasing on average with distance from the Galactic Center.[40][48][49][50][51]

Galaxies

In November 2022, astronomers, using the Hubble Space Telescope, discovered one of the most metal-poor galaxies known. This nearby dwarf galaxy, 20 million ly away and 1,200 ly across, is named HIPASS J1131–31 (nicknamed the "Peekaboo" Galaxy).[52][53] According to one of the astronomers, "Due to Peekaboo's proximity to us, we can conduct detailed observations, opening up possibilities of seeing an environment resembling the early universe in unprecedented detail."[54]

See also

References

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  3. ↑ White, S.D.M.; Rees, M.J. (1978). "Core condensation in heavy halos - a two-stage theory for galaxy formation and clustering". Monthly Notices of the Royal Astronomical Society 183 (3): 341–358. doi:10.1093/mnras/183.3.341. Bibcode: 1978MNRAS.183..341W. 
  4. ↑ Puget, J.L.; Heyvaerts, J. (1980). "Population III stars and the shape of the cosmological black body radiation". Astronomy and Astrophysics 83 (3): L10–L12. Bibcode: 1980A&A....83L..10P. 
  5. ↑ Asplund, Martin; Grevesse, Nicolas; Sauval, A. Jacques; Scott, Pat (2009). "The chemical composition of the Sun". Annual Review of Astronomy & Astrophysics 47 (1): 481–522. doi:10.1146/annurev.astro.46.060407.145222. Bibcode: 2009ARA&A..47..481A. 
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  27. ↑ Claria, Juan J.; Piatti, Andres E.; Lapasset, Emilio (May 1994). "A revised effective-temperature calibration for the DDO photometric system". Publications of the Astronomical Society of the Pacific 106: 436. doi:10.1086/133398. ISSN 0004-6280. Bibcode: 1994PASP..106..436C. 
  28. ↑ James, K.A. (May 1975). "Cyanogen strengths, luminosities, and kinematics of K giant stars". The Astrophysical Journal Supplement Series 29: 161. doi:10.1086/190339. ISSN 0067-0049. Bibcode: 1975ApJS...29..161J. 
  29. ↑ Wang, Ji. "Planet-metallicity correlation - the rich get richer". Caltech. http://www.astro.caltech.edu/~jwang/Project4.html. 
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  32. ↑ Sanders, Ray (9 April 2012). "When stellar metallicity sparks planet formation". Astrobiology Magazine. http://www.astrobio.net/news-exclusive/when-stellar-metallicity-sparks-planet-formation/. 
  33. ↑ "The G star problem". IAU Symposium 228. 228. pp. 509–511.  [citation not found]
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  34. ↑ Kewley, L.J.; Dopita, M.A. (September 2002). "Using strong lines to estimate abundances in extragalactic HII regions and starburst galaxies". The Astrophysical Journal Supplement Series 142 (1): 35–52. doi:10.1086/341326. ISSN 0067-0049. Bibcode: 2002ApJS..142...35K. 
  35. ↑ Nagao, T.; Maiolino, R.; Marconi, A. (2006-09-12). "Gas metallicity diagnostics in star-forming galaxies". Astronomy & Astrophysics 459 (1): 85–101. doi:10.1051/0004-6361:20065216. ISSN 0004-6361. Bibcode: 2006A&A...459...85N. 
  36. ↑ Peimbert, Manuel (December 1967). "Temperature determinations of HII regions". The Astrophysical Journal 150: 825. doi:10.1086/149385. ISSN 0004-637X. Bibcode: 1967ApJ...150..825P. 
  37. ↑ Pagel, B.E.J. (1986). "Nebulae and abundances in galaxies". Publications of the Astronomical Society of the Pacific 98 (608): 1009. doi:10.1086/131863. ISSN 1538-3873. Bibcode: 1986PASP...98.1009P. http://stacks.iop.org/1538-3873/98/i=608/a=1009. 
  38. ↑ Henry, R.B.C.; Worthey, Guy (August 1999). "The distribution of heavy elements in spiral and elliptical galaxies". Publications of the Astronomical Society of the Pacific 111 (762): 919–945. doi:10.1086/316403. ISSN 0004-6280. Bibcode: 1999PASP..111..919H. 
  39. ↑ Kobulnicky, Henry A.; Kennicutt, Robert C. Jr.; Pizagno, James L. (April 1999). "On measuring nebular chemical abundances in distant galaxies using global emission‐line spectra". The Astrophysical Journal 514 (2): 544–557. doi:10.1086/306987. ISSN 0004-637X. Bibcode: 1999ApJ...514..544K. 
  40. ↑ 40.0 40.1 Grazyna, Stasinska (2004). "Abundance determinations in HII regions and planetary nebulae". Cosmochemistry: The melting pot of the elements. Cambridge Contemporary Astrophysics. Cambridge University Press. pp. 115–170. Bibcode: 2002astro.ph..7500S. 
  41. ↑ Peimbert, Antonio; Peimbert, Manuel; Ruiz, Maria Teresa (December 2005). "Chemical composition of two HII regions in NGC 6822 based on VLT spectroscopy". The Astrophysical Journal 634 (2): 1056–1066. doi:10.1086/444557. ISSN 0004-637X. Bibcode: 2005ApJ...634.1056P. 
  42. ↑ Pagel, B.E.J.; Edmunds, M.G.; Blackwell, D.E.; Chun, M.S.; Smith, G. (1979-11-01). "On the composition of HII regions in southern galaxies – I. NGC 300 and 1365". Monthly Notices of the Royal Astronomical Society 189 (1): 95–113. doi:10.1093/mnras/189.1.95. ISSN 0035-8711. Bibcode: 1979MNRAS.189...95P. 
  43. ↑ Dopita, M.A.; Evans, I.N. (August 1986). "Theoretical models for HII regions. II - The extragalactic HII region abundance sequence" (in en). The Astrophysical Journal 307: 431. doi:10.1086/164432. ISSN 0004-637X. Bibcode: 1986ApJ...307..431D. 
  44. ↑ McGaugh, Stacy S. (October 1991). "HII region abundances - Model oxygen line ratios". The Astrophysical Journal 380: 140. doi:10.1086/170569. ISSN 0004-637X. Bibcode: 1991ApJ...380..140M. 
  45. ↑ Pilyugin, L.S. (April 2001). "On the oxygen abundance determination in HII regions". Astronomy & Astrophysics 369 (2): 594–604. doi:10.1051/0004-6361:20010079. ISSN 0004-6361. Bibcode: 2001A&A...369..594P. https://www.aanda.org/articles/aa/ps/2001/14/aa10396.ps.gz. 
  46. ↑ Kobulnicky, Henry A.; Zaritsky, Dennis (1999-01-20). "Chemical Properties of Star‐forming Emission‐Line Galaxies atz=0.1–0.5". The Astrophysical Journal 511 (1): 118–135. doi:10.1086/306673. ISSN 0004-637X. Bibcode: 1999ApJ...511..118K. 
  47. ↑ Diaz, A.I.; Perez-Montero, E. (2000-02-11). "An empirical calibration of nebular abundances based on the sulphur emission lines". Monthly Notices of the Royal Astronomical Society 312 (1): 130–138. doi:10.1046/j.1365-8711.2000.03117.x. ISSN 0035-8711. Bibcode: 2000MNRAS.312..130D. 
  48. ↑ Shaver, P.A.; McGee, R.X.; Newton, L.M.; Danks, A.C.; Pottasch, S.R. (1983-09-01). "The galactic abundance gradient". Monthly Notices of the Royal Astronomical Society 204 (1): 53–112. doi:10.1093/mnras/204.1.53. ISSN 0035-8711. Bibcode: 1983MNRAS.204...53S. 
  49. ↑ Afflerbach, A.; Churchwell, E.; Werner, M. W. (1997-03-20). "Galactic abundance gradients from infrared fine‐structure lines in compact HII regions". The Astrophysical Journal 478 (1): 190–205. doi:10.1086/303771. ISSN 0004-637X. Bibcode: 1997ApJ...478..190A. 
  50. ↑ Pagel, J.; Bernard, E. (1997). Nucleosynthesis and Chemical Evolution of Galaxies. Cambridge University Press. p. 392. ISBN 978-0-521-55061-1. Bibcode: 1997nceg.book.....P. 
  51. ↑ Balser, Dana S.; Rood, Robert T.; Bania, T.M.; Anderson, L.D. (2011-08-10). "HII region metallicity distribution in the Milky Way disk". The Astrophysical Journal 738 (1): 27. doi:10.1088/0004-637X/738/1/27. ISSN 0004-637X. Bibcode: 2011ApJ...738...27B. 
  52. ↑ Karachentsev, J.D. (12 November 2022). "Peekaboo: The extremely metal poor dwarf galaxy HIPASS J1131-31". Monthly Notices of the Royal Astronomical Society 518 (4): 5893–5903. doi:10.1093/mnras/stac3284. https://academic.oup.com/mnras/article-abstract/518/4/5893/6825465. Retrieved 17 December 2022. 
  53. ↑ Villard, Ray (6 December 2022). "Peekaboo! A tiny, hidden galaxy provides a peek into the past - tucked away in a local pocket of dark matter, a late-blooming dwarf galaxy looks like it belongs in the early universe". Hubblesite.org (Press release). NASA. Retrieved 18 December 2022.
  54. ↑ Parks, Jake (16 December 2022). "Hubble spots a nearby galaxy that looks like it belongs in the early universe - The extremely metal-poor galaxy, nicknamed 'Peekaboo', relatively recently emerged from behind a fast-moving star". https://astronomy.com/news/2022/12/hubble-spots-a-nearby-galaxy-that-looks-like-it-belongs-in-the-early-universe. 

Further reading



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