Atomic Hydrogen in Distant Galaxies∗ Nissim Kanekar, Aditya Chowdhury, and Jayaram N. Chengalur Atomic hydrogen is a fundamental constituent of galaxies, and is the primary fuel for star formation. Understanding how the atomic hydrogen content of galaxies evolves with time is thus critical to understanding galaxy evolution. In this ar- ticle, we discuss how one measures the atomic hydrogen con- tent of distant galaxies, and what these measurements tell us Nissim Kanekar, Aditya about the evolution of our Universe. Chowdhury, and Jayaram N. Chengalur are astronomers at the National Centre for Radio 1. Stars and the Interstellar Medium Astrophysics, Tata Institute of Fundamental Research, Pune. The Milky Way contains about a hundred billion stars. Although this is a staggeringly large number, the stars take up only a tiny fraction of the volume of the Milky Way, with vast spaces in be- tween them. These spaces are occupied by gas and dust, making up what is known as the interstellar medium (ISM). The ISM con- sists mostly of hydrogen gas, in neutral atomic, neutral molecular, or ionized forms. The mass of dust in the ISM is very small com- pared to the mass in gas, but the dust plays an important role in the heating and cooling of the ISM. The evolution of galaxies like the Milky Way is driven by com- plex interactions between the stars and the ISM. A simplified se- quence of the critical processes is: (i) warm ionized hydrogen flows into a galaxy from its surroundings, i.e., from the circum- galactic medium (CGM), (ii) the gas settles down into the disk of the galaxy, cools, and forms neutral atomic hydrogen (Hi), Keywords (iii) small regions of the neutral atomic hydrogen cool further to Galaxy evolution, High-redshift galaxies, atomic hydrogen, form molecular hydrogen (H2), (iv) the dense molecular clouds HI 21cm line, GMRT, Govind collapse and fragment under their own gravity to form stars, (v) at Swarup. the end of their lives, the more massive stars blow up as super- ∗Vol.26, No.7, DOI: https://doi.org/10.1007/s12045-021-1192-2 RESONANCE | July 2021 919 novae, throwing higher elements (referred to as “metals”, e.g., carbon, oxygen, zinc, iron, etc.) into the ISM, and (vi) the metals 1Metal atoms typically have give rise to further cooling of the gas,1 causing further star forma- upper energy levels to which tion, and so on. This sequence continues in a galaxy until most their electrons can be excited of its hydrogen has been consumed in star formation, after which via collisions. The metal atom subsequently returns to a lower the galaxy stops forming new stars. energy state with the emission Thus, one can think of the neutral atomic and molecular hydrogen of a photon. These emitted pho- fuel tons typically carry away en- as the for star-formation in a galaxy. Stars can form in a ergy from the ISM, and thus ef- galaxy only as long as it contains atomic and molecular hydrogen. fectively cool it. If the molecular hydrogen runs out, star-formation would cease. Conversely, if the existing atomic hydrogen runs out, there would be no raw material for the formation of H2, and stars can form For the atomic gas, the only as long as the H2 reservoir lasts. Hi and H2 thus play key sole way of measuring roles in regulating how stars form, and when the star-formation i the H content of process ends. galaxies is via the Hi hyperfine transition at Today’s galaxies typically contain a far larger mass in stars than radio wavelengths, the i in gas: for example, the Milky Way contains about 60 billion so- famous “H 21cm” i spectral line. lar masses in stars, but only around 8 billion solar masses in H and around 2.5 billion solar masses in H2. Stars are thus dominant over gas in the bigger galaxies today. But was this the situation 2Molecular gas in the Uni- ten billion years ago? What were galaxies like when they were verse is primarily H2. How- young and reckless? Understanding galaxy evolution requires ever, H2 is a symmetric, very us to understand the evolution of both the stars and the neutral light molecule, with no strong spectral lines that can be de- atomic and molecular gas in galaxies. tected from the cold clouds that The stellar properties of galaxies are best studied in the ultravi- make up the bulk of the molec- ular ISM. The strong mm-wave olet, optical, and near-infrared wavelengths, where the stars are lines of the second-most abun- very bright. In the case of the molecular gas, the strong spectral dant molecule, CO, are hence lines of the carbon monoxide molecule at millimetre wavelengths used as a “trace” of the molec- provide the best tool to probe the molecular content of galaxies.2 ular gas. And, for the atomic gas, the sole way of measuring the Hi content of galaxies is via the Hi hyperfine transition at radio wavelengths, the famous “Hi 21cm” spectral line. 920 RESONANCE | July 2021 2. TheHi 21cm Spectral Line The science of radio astronomy was born to engineering parents, Karl Jansky and Grote Reber, in the 1930s in the USA. Both Jansky and Reber detected broad-band emission signals from the Milky Way, especially from the centre of our galaxy. In war-torn Holland in the early 1940s, the great Dutch astronomer Jan Oort came across Reber’s work, and realized the importance of find- ing a spectral line at radio wavelengths. The beauty of spectral lines is that each line arises from emission or absorption at a very specific wavelength (or a narrow wavelength range around a spe- cific wavelength). If the source of the spectral line (e.g., a star or an ISM cloud) is moving towards us, the line would be seen to be shifted to shorter wavelengths due to the Doppler effect; 3The scattering of light from conversely, the line would shift to longer wavelengths for sources dust particles is ∝ λ−4, i.e., moving away from us, for the same reason. For optical lines, the is much higher at short wave- lengths than at long wave- former shift is towards the blue end of the spectrum, and is re- lengths, if the size of the dust ferred to as a blueshift, while the latter shift is towards the red grains are smaller than the end of the spectrum and is referred to as a redshift. Crucially, the wavelength (“Rayleigh scatter- Doppler effect allows us to measure the velocity of the source of ing”). The observed depen- dence of the scattering on wave- the line! length has been used to show Oort had mapped the rotation of the Milky Way using optical that dust grains in the ISM have sizes of roughly 0.01 − 0.2µm. spectral lines and had found that optical lines are difficult to de- This is why the scattering is tect in directions containing large amounts of dust (e.g., towards so strong at ultraviolet wave- the centre of the Milky Way) because the dust absorbs or scat- lengths, but negligible at the ters the optical radiation. He had earlier argued for the presence much longer radio wavelengths. of large gas clouds in the Milky Way and understood that a ra- dio spectral line could be used to trace these gas clouds. And, because radio spectral lines3 are at centimetre or metre wave- lengths, much much larger than the typical size of the dust grains van de Hulst calculated (roughly, 0.01 − 0.2 microns), radio radiation would pass unhin- the strengths of different dered through any dust along the path. Radio spectral lines would possible radio lines and predicted that it might be thus be superb tracers of the structure of the Milky Way! possible to detect a In the spring of 1944, Oort shared his enthusiasm to find such hyperfine transition in the hydrogen atom, at a a radio spectral line with Hendrik van de Hulst, a student at the wavelength of roughly University of Utrecht. van de Hulst calculated the strengths of 21 cm (the Hi 21cm line). RESONANCE | July 2021 921 different possible radio lines, and predicted that it might be pos- sible to detect a hyperfine transition in the hydrogen atom, at a wavelength of roughly 21 cm (the Hi 21cm line) (see Box 1). Box 1. The Hi 21cm line Figure A. A hydrogen atom jumps from the state with parallel proton and electron spins to the state with anti-parallel spins, with the emission of a photon of wavelength 21 cm. A hydrogen atom consists of a proton and an electron. Each particle possesses a “spin”, and hence, a magnetic moment. In a hydrogen atom, there are two possible quantum states arising from the orientation of the electron spin relative to the proton spin, one in which the spins are parallel and the other in which the spins are anti-parallel. Due to the interaction of the magnetic moments of the two particles, the two states (called the “hyperfine” states) have a small difference in their energy. The hydrogen atom can spontaneously “jump” from the state with parallel spins to the state with anti-parallel spins, with the emission of a photon with energy equal to the energy difference between the states. This energy difference is tiny compared to the difference between, for example, the first and second Bohr orbits. Since the energy of the photon is inversely proportional to its wavelength, the hyperfine photon arises at a very long wavelength, roughly 21 cm.