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Hydrogen is the most abundant elementin the interstellar medium(ISM), but the symmetric H$_2$ molecule has no permanent dipole momentand hence does not emit a detectable spectral line at radiofrequencies. Neutral hydrogen (HI) atoms are abundant andubiquitous in low-density regions of the ISM. They are detectable inthe $lambda approx 21$ cm ($nu_{10} = 1420.405751...$ MHz) hyperfineline. Two energy levels result from the magneticinteraction between the quantized electron and proton spins. Whenthe relative spins change from parallel to antiparallel, a photon isemitted.
One $lambda = 21$ cm photon isemitted when the spins flip from parallel to antiparallel.Imagecredit
The line center frequency is
$$bbox[border:3px blue solid,7pt]{nu_{10} = {8 over 3} g_{rm I}biggl( {m_{rm e} over m_{rm p} }biggr) alpha^2 (R_{rm M} c)approx1420.405751{rm ~MHz}}rlap{quad rm {(7E1)}}$$ where$g_{rm I} approx 5.58569$ is the nuclear$g$-factorfor aproton,$alpha equiv e^2 / (hbar c) approx 1 / 137.036$ is thedimensionless fine-structureconstant,and $R_{rm M}c$ is the hydrogen Rydberg frequency (Eq. 7A2).
$$nu_{10} approx { 8 over 3} cdot5.58569 cdot biggl( { 1 over 1836.12} biggr) biggl( { 1 / 137.036}biggr)^2 cdot 3.28984 times 10^{15} {rm ~Hz~}
biggl(1 + {1 over 1836.12}biggr)^{-1}$$ $$nu_{10} approx 1420.4{rm ~MHz}$$
By analogy with the emissioncoefficient of radiation by an electricdipole $$A_{rm UL} approx { 64 pi^4 over 3 h c^3} nu_{rm UL}^3vert mu_{rm UL} vert^2~,$$ the emission coefficient of thismagneticdipole is $$A_{rm UL} approx { 64 pi^4 over 3 h c^3} nu_{rm UL}^3vert mu_{rm 10}^* vert^2~,$$ where $mu_{rm 10}^*$ is the meanmagnetic dipole moment for HI in the ground electronic state($n = 1$). The magnitude $vert mu_{rm 10}vert$ equals the Bohrmagneton, the intrinsic dipole moment of an electron. Electrons have spinangular momentum $L = hbar /2$, classical radius $r_{rm e} = e^2 /(m_{rm e}c^2)$, and charge $e$, so
$$vert mu_{10}^* vert = {hbar over 2}{e over m_{rm e} c}approx {6.63 times 10^{-27} {rm ~erg~s}/(2pi) over 2} cdot { 4.8times10^{-10} {rm ~statcoul~} over 9.11 times 10^{-28} {rm ~g~}cdot 3 times10^{10} {rm ~cm~s}^{-1} }$$ $$vert mu_{10}^* vert approx 9.27times 10^{-21} {rm~erg~Gauss}^{-1} $$ Thus the emission coefficient of the 21 cm line is
$$A_{10} approx { 64 pi^4 (1.42 times 10^9 {rm ~Hz})^3 over 3cdot 6.63 times 10^{-27} {rm ~erg~s~} (3 times 10^{10} {rm~cm~s}^{-1})^3 } (9.27 times 10^{-21} {rm ~erg~Gauss}^{-1})^2$$$$bbox[border:3px blue solid,7pt]{A_{10} approx2.85 times 10^{-15} {rm ~s}^{-1}}rlap{quad rm {(7E2)}}$$ That is, the radiative half-lifeof this transition is about
$$tau_{1/2} = A_{10}^{-1} approx3.5 times 10^{14} {rm ~s~} approx 11 {rm ~million~years}$$ Such alow emission coefficient implies an extremely low critical density(defined by Eq. 7D10) $n^* ll 1 {rm ~cm}^{-3}$, so collisions caneasily maintain this transition in LTE throughout the diffuseinterstellar medium of a normal galaxy.
Regardlessof whether the HI is in LTE or not, we can define the HIspin temperature $T_{rm s}$ (the HI analog of the molecularexcitation temperature $T_{rm x}$ defined by Equation 7B8) by
$$bbox[border:3px blue solid,7pt]{{N_1 over N_0} equiv {g_1 overg_0} exp biggl( - {hnu_{10} over k T_{rm s} } biggr)}rlap{quad rm {(7E3)}}~,$$ where the statistical weightsof the upper and lower spin states are $g_1 = 3$ and $g_0 = 1$,respectively. Note that
$$ {h nu_{10} over k T_{rm s}} approx { 6.63 times10^{-27} {rm ~erg~s~} cdot 1.42 times 10^9 {rm ~Hz} over 1.38times 10^{-16} {rm ~erg~K}^{-1} cdot 150 {rm ~K} } approx 5 times10^{-4} ll 1$$ is very small for gas in LTE at $T approx T_{rm s}approx 150$ K, so in the ISM
$$ {N_1 over N_0} approx{g_1 over g_0} = 3 qquad {rm and} qquad N_{rm H} = N_0 + N_1approx 4 N_0$$
Inserting these weights into Equation7B7 for the opacity coefficient of the $lambda = 21 {rm ~cm}$line gives $$ kappa_nu = {c^2 over 8 pi nu_{10}^2 } {g_1 overg_0} N_0A_{10} biggl[ 1 - exp biggl( - { h nu_{10} over k T_{rm s}}biggr) biggr] phi(nu)$$ $$kappa_nu approx { c^2 over 8 pinu_{10}^2 } cdot 3 cdot {N_{rm H} over 4} A_{10} biggl( { hnu_{10} over k T_{rm s}} biggr) phi(nu)$$
$$bbox[border:3px bluesolid,7pt]{kappa_nu approx { 3 c^2 over 32pi} {A_{10} N_{rm H} overnu_{10} } {h over k T_{rm s}} phi (nu) ~}rlap{quad rm{(7E4)}}$$ where $N_{rm H}$ is the numberofneutral hydrogen atoms percm$^{3}$. The neutral hydrogen column densityalong any line-of-sight is defined as
$$bbox[border:3px blue solid,7pt]{ eta_{rm H} equiv int_{rm los}N_{rm H} (s) d s}rlap{quad rm {(7E5)}}$$ The total opacity $tau$ of isothermal HI is proportional to the columndensity. If $tau ll 1$, then the integrated HIemission-line brightness $T_{rm b}$ isproportional to the column density of HIand is independent of the spin temperature $T_{rm s}$ because $T_{rmb} approx T_{rm s} tau$ and $tau propto T_{rm s}^{-1}$ in theradiolimit $h nu_{10} / (k T_{rm s}) ll 1$. Thus $eta_{rm H}$ canbe determined directly from the integrated line brightness when $taull 1$. Inastronomicallyconvenient units it can be written as
$$bbox[border:3px blue solid,7pt]{biggl( { eta_{rm H} over {rmcm}^{-2} }biggr) approx1.82 times 10^{18} int biggl[ { T_{rm b} (v) over {rm K}} biggr]d biggl( { v over {rm km~s}^{-1} } biggr)}rlap{quad rm{(7E6)}}$$ where $T_{rm b}$ is theobserved 21 cm line brightness temperatureat radial velocity $v$ and the velocity integration extends over theentire 21 cm line profile. Note that absorptionby HIin front of a continuum source with continuum brightness temperagure$> T_{rm s}$, on theother hand, is weighted in favor of colder gas.
The HI absorption and emissionspectra toward the source 1714-397 (Dickey, J. M. et al. 1983, ApJS,53, 591).
The equilibrium temperature of coolinterstellar HI isdetermined by the balance of heating and cooling. The primary heatsources are cosmic rays and ionizing photons from hot stars. The maincoolant in the cool ISM is radiation from the fine-structure line ofsingly ionizedcarbon, CII, at $lambda = 157.7,mu$m. This line is strongonly when the temperature is at least
$$k T approx h nu = {h c overlambda}~,$$ so the cooling rate increases exponentially above
$$ Tapprox {h c over k lambda} approx {6.63 times 10^{-27} {rm~erg~s} cdot 3 times 10^{10} {rm ~cm~s}^{-1} over 1.38 times10^{-16} {rm ~erg~K}^{-1} cdot 157.7 times 10^{-4} {rm ~cm}}approx 91~K~.$$ The actual kinetic temperature of HI in ourGalaxy can be estimated from the HI line brightnesstemperatures in directions where the line is optically thick ($tau gg1$) and the brightness temperature approaches the excitationtemperature, which is close to the kinetic temperature in LTE. Manylines-of-sight near the galactic plane have brightness temperatures ashigh as 100–150 K, values consistent with the temperature-dependentcooling rate.
Galactic HI
Neutral hydrogen gas in the disk ofour Galaxy moves in nearly circular orbits around the galacticcenter. Radial velocities $V_{rm r}$ measured from the Dopplershifts of HI $lambda = 21$ cm emission lines encode information aboutthe kinematicdistances $d$ of HI clouds, and the spectra of HI absorption infront of continuum sources can be used to constrain their distancesalso. HI is optically thin except in a few regions near thegalactic plane, so the distribution of hydrogen maps out thelarge-scale structure of the whole Galaxy, most of which is hidden bydust at visible wavelengths.
The figure below shows a plan view ofthe galactic disk. The Sun ($odot$) lies in the disk and movesin a circular orbit around the galactic center. The distance tothe galactic center $R_odot = 8.0 pm 0.5$ kpc and the Sun's orbitalspeed $omega_odot R_odot approx 220$ km s$^{-1}$ have been measuredby a variety of means (Reid, M. J. 1993, ARAA, 31, 345). All HIclouds at galactocentric distance $R$ are assumed to be in circularorbits with angular velocity $omega(R)$, where $omega(R)$ is amonotonically decreasing function of $R$. For cloud 1 atgalactocentric azimuth $theta$ on the line of sight at galacticlongitude $l$, the observed radial velocity $V_{rm r}$ relative to theSun is given by
$$V_{rm r} = omega R cos[pi/2 - (l + theta)] - omega_odotR_odot cos(pi/2 - l)$$
Using the trigonometric identities $cos[pi/2 - (l + theta)] = sin(l+ theta)$ and $sin(l+theta) = sintheta cos l + costheta sin l$we obtain
$$V_{rm r} = omega R (sintheta cos l + costheta sin l) -omega_odot R_odot sin l$$ $$V_{rm r} = R_odot (omega-omega_odot)sin l$$
To apply this equation, we need to determine the rotation curve$Romega(R)$. The maximum radial velocity on the line of sight atlongitude $l$ is called the 'terminal velocity' $V_{rm T}$. Since$omega$ decreases with $R$, this velocity occurs at the minimum $R =R_{rm min} = R_odot sin l$ where the orbit is tangent to the line ofsight:
$$V_{rm T} = R_odot [omega(R_{rm min})-omega_odot]sin l $$ Wecan determine the rotation curve from measurements of $V_{rm T}$spanning a wide range of $l$ and thus of $R_{rm min}$.
$$R_{rm min} = R_odot sin l = 8.0 {rm ~kpc} cdot 0.5 = 4.0 {rm~kpc}$$ $$ V_{rm T} =R_odot [omega(R_{rm min}) - omega_odot]sin l$$ $$ V_{rm T} =R_{rm min} omega(R_{rm min}) - R_odot omega_odot sin l$$ $$R_{rm min} omega(R_{rm min}) = V_{rm T} + R_odot omega_odot sinl$$ $$R_{rm min } omega(R_{rm min}) = 130 {rm ~km~s}^{-1} + 220{rm ~km~s}^{-1} times 0.5
= 240 {rm ~km~s}^{-1}$$
Note that, for $vert l vert <pi /2$, there is a distance ambiguity: clouds 1 and 2 both have thesame radial velocities but different distances $d$. There is nodistance ambiguity for $vert l vert > pi/2$.
In the simplest realistic model for the galactic disk, the Sunand all HI clouds are in circular orbits about the Galactic Center, andthe angular orbital velocity $omega$ is a monotonically decreasingfunction of the orbital radius $R$. The distance of the GalacticCenter is $R_odot = 8.0 pm 0.5$ kpc, and the Sun's orbital speedis $omega_odot R_odot approx 220$ km s$^{-1}$. For $vert lvert < pi /2$, two HI clouds (1 and 2) can have the same radialvelocity but be at different distances from the Sun.
Thisimage shows a galactic 'superbubble' in HI (green) and HII (purple)about 7 kpc distant and 3 kpc in height. Stellar winds andsupernovae in young star clusters blow these bubbles. Images ofHI away from the galactic plane are easily contaminated by sideloberesponses to the strong and widespread HI emission from the planeitself. The low sidelobe levels of the clear-aperture GBT makesuch HI images possible. Imagecredit
HI in ExternalGalaxies
The 1420 MHz HI line is an extremelyuseful tool for studying gas in the ISM of external galaxies andtracing the large-scale distribution of galaxies in the universebecause HI is detectable in most spiral galaxies and in some ellipticalgalaxies.
The observed center frequency of theHI line can be used to measure the radial velocity $V_{rm r}$ of agalaxy. The radial velocity of a galaxy is the sum of the recessionvelocity caused by the uniform Hubble expansion of the universe and the'peculiar' velocity of the galaxy. The radial component of thepeculiar velocity reflects motions caused by gravitational interactionswith nearby galaxies and is typically $sim 200$ km s$^{-1}$ inmagnitude. The Hubble velocity is proportional to distance fromthe Earth, and the Hubbleconstant of proportionality has been measured as $H_0 approx72$ km s$^{-1}$ Mpc$^{-1}$. If the radial velocity issignificantly larger than the radial component of the peculiarvelocity, the observed HI frequency can be used to estimate the Hubbledistance $D approx v_{rm r} / H_0$ to a galaxy.
Example: Use the HI emission-lineprofile below of the galaxy UGC 11707to estimate its Hubble distance
$D approx v_{rm r} / H_0$ if $H_0 approx 72$ km s$^{-1}$ Mpc$^{-1}$.If the radial velocity$v_{rm r} ll c$, then the nonrelativistic Doppler formula can be usedto calculate $v_{rm r}$ from the observed line frequency $nu$:
$${v_{rm r} over c}approx {nu_0 - nu over nu_0}~,$$ where $nu_0 approx 1420.4$ MHzis the rest-frame frequency. This equationyields what is known as the radiovelocity because radioastronomers measure frequencies, not wavelengths. Optical astronomersmeasure wavelengths, not frequencies, so theoptical velocityis defined by $${v_{rm r} over c} approx {lambda - lambda_0 overlambda_0}~.$$ Beware of this 'gotcha': the optical and radiovelocitiesare not exactly equal. Occasionally an observer confuses them, failsto center the observing passband on the correct frequency, and ends upwith only part of the HI spectrum of a galaxy.
Since $lambda = 21$ cm is such a longwavelength, many galaxies areunresolved by single-dish radio telescopes. For example, the half-powerbeamwidth of the 100 m GBT is about 9 arcmin at $lambda = 21$ cm. Thusa single pointing is sufficient to obtain a spectral line representingall of the HI in any but the nearest galaxies.
This integrated HI spectrum of UGC 11707 obtained with the 140-foottelescope (beamwidth $approx 20$ arcmin) shows the typical two-hornedprofile of a spiral galaxy.
For UGC 11707, the line centerfrequency is $nu approx1416.2$ MHz, so $$v_{rm r} approx c biggl( 1 - { nu over nu_0}biggr) approx 3 times 10^5 {rm ~km~s}^{-1} biggl( 1 - {1416.2 {rm~MHz} over 1420.4 {rm ~MHz}} biggr) approx 890 {rm ~km~s}^{-1}$$$$D approx {v_{rm r} over H_0} = {890 {rm ~km~s}^{-1} over 72 {rm~km~s}^{-1} {rm ~Mpc}^{-1} } = 12.4 {rm ~Mpc}$$
If the HI emission from a galaxy isoptically thin, then the integrated line flux is proportional to themass of HI in the galaxy, independent of the unknown HItemperature. It is a straightforward excercise to derive fromEquation 7E6 the relation $$bbox[border:3px blue solid,7pt]{biggl ({M_{rm H} over M_odot}biggr) approx 2.36 times10^5 biggl( { D over {rm Mpc} } biggr)^2 int biggl[ { S(v) over{rm Jy}} biggr] biggl( {d v over {rm km~s}^{-1} }biggr)}rlap{quad rm {(7E7)}}$$ for thetotal HImass $M_{rm H}$ of a galaxy. Theintegral $int S(v) d v$ over theline iscalled theline flux and is usually expressed in units of Jy km s$^{-1}$.
Example: What is the HImass of UGC 11707?
We assume $tau ll 1$. Thesingle-dish profile of UGC 11707 above indicates a line flux
$$int S(v), dv approx 0.35 {rm ~Jy~} times 200 {rm ~km~s}^{-1} approx 70 {rm~Jy~km~s}^{-1}$$ so $$biggl( {M_{rm H} over M_odot} biggr) approx2.36 times 10^5 cdot (12.4)^2 cdot 70 approx 2.5 times 10^9$$Small statistical corrections for nonzero $tau$ can be made fromknowledge about the expected opacity as a function of disk inclination,galaxy mass, morphological type, etc.
A well-resolved HI image of a galaxyyields the totalmass $M(r)$ enclosed within radius $r$ of the center if thegas orbits in circular orbits. Circular orbitsyield a characteristic radial velocity field, illustrated by nearbyspiral galaxy M33.
The HI radial velocity field of thenearby spiral galaxy M33 is shown here by colors corresponding toDoppler redshifts and blueshifts relative to the center of mass;brightness in this image is proportional to HI column density. Imagecredit
If the mass distribution of everygalaxy were spherically symmetric, the gravitational force at radius$r$ would equal the gravitational force of the enclosed mass$M(r)$. This is not a bad approximation, even for diskgalaxies. Thus for gas in circular orbits,
$${ G Mover r^2} = {v^2 over r}~,$$where $M$ is the mass enclosed withinthe sphere of radius $r$ and $v$ is the orbital velocity at radius $r$,so
$$v^2 = { G M over r}~.$$ Notethat the velocity $v$ is the full tangential velocity, not just theradial component $v_{rm r}(r)$ that contributes to the observableDoppler shift:
$$v(r)= {v_{rm r}(r) over sin i}~,$$ where $i$ is the inclination anglebetween the galaxy disk and the line-of-sight. The inclination angle ofa thin circular disk can be estimated from the axial ratio
$$cos i ={theta_{rm m} over theta_{rm M}}~,$$ where $theta_{rm m}$ and$theta_{rm M}$ are the minor- and major-axis angular diameters,respectively. Converting from cgs to astronomically convenient unitsyields
$$biggl[ biggl( { v over {rm cm~s}^{-1} } biggr) biggl( { 10^5{rm ~cm~s}^{-1} over {rm km~s}^{-1} } biggr) biggr]^2 = $$
$$biggl[6.67 times 10^{-8} {rm ~dyne~cm}^2 {rm ~g}^{-2} cdot biggl( { Mover {rm g} } biggr) biggl( { 2 times 10^{33} {rm ~g} overM_odot} biggr) biggr] biggl[ biggl( { r over {rm cm}} biggr)biggl( { 3.09 times 10^{21} {rm ~cm} over {rm kpc} } biggr)biggr]^{-1}$$
$$10^{10} biggl( { v over {rm km~s}^{-1} } biggr)^2= biggl[ 6.67 times 10^{-8} cdot 2 times 10^{33} biggl( { M overM_odot} biggr) biggr] biggl[ 3.09 times 10^{21} biggl( { r over{rm kpc}} biggr) biggr]^{-1}$$ and we obtain the total galaxy massinside radius $r$ in unitsof the solar mass:
$$bbox[border:3px blue solid,7pt]{ biggl( { M over M_odot} biggr)approx 2.3 times 10^5biggl( { v over {rm km~s}^{-1} } biggr)^2 biggl( { r over {rmkpc} } biggr) approx 2.3 times 10^5 biggl[ { (v_{rm r} / sin i)over {rm km~s}^{-1} } biggr]^2 biggl( { r over {rm kpc} } biggr)}rlap{quad rm {(7E8)}}$$
Example: What is the total mass ofUGC 11707?
HI images of UGC 11707 (Swaters, R.A. et al. 2002, A&A, 390, 829). The contours in panels (a)and (c) outline the integrated HI brightness distribution. Panel(b) shows contours of constant velocity separated by 20 km s$^{-1}$ andthe darker shading indicates approaching gas. Panel (d) is aposition-velocity diagram, panel (e) is the radial HI column-densityprofile, and panel (f) displays the integrated HI spectrum.
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$$ v_{rm r} approx {Delta v_{rm r}over 2} approx { (1000 {rm~km~s}^{-1} - 800 {rm ~km~s}^{-1}) over 2} approx 100 {rm~km~s}^{-1}$$
$$ cos i approx { {rm minor~axis} over {rmmajor~axis} } approx { 0.73 times 10^{-3} {rm ~rad} over 2.0 times10^{-3} {rm ~rad} } approx 0.365 qquad {rm so} qquad sin iapprox 0.93$$
$$r approx theta_{1/2} D approx 10^{-3} {rm ~rad~}cdot 12.4{rm ~Mpc} approx 12.4 {rm ~kpc}$$ so
$$biggl( { M overM_odot} biggr) approx 2.3 times 10^5 cdot (100 / 0.93) ^2 cdot12.4 = 3.3 times 10^{10}~.$$ UGC 11707 is a relatively low-mass spiralgalaxy.
This 'total' mass is really only themass inside the radius sampledby detectable HI. Even though HI extends beyond mostother tracers such as molecular gas or stars, it is clear fromplots of HI rotation velocities versus radius that not allof the mass is being sampled, because we don't see the Keplerianrelation $v_{rm r} proptor^{-1/2}$ which indicates that all of the mass is enclosed withinradius $r$. Most rotationcurves, one-dimensional position-velocity diagrams along themajor axis, are flat atlarge $r$, suggesting that theenclosed mass $M propto r$ as far as we can see HI. The largetotal masses implied by HI rotation curves provided some of theearliest evidence for the existence of cold dark matter in galaxies.
This postershows VLA HI images of THINGS (The HI Survey of NearbyGalaxies) galaxies at constant linear scale and linear resolution. Imagecredit
Ahigh-resolution HI image of M81 made with the VLA for the THINGS(de Blok et al. 2007, astro-ph/0407103)survey compared with Spitzer mid-infraredemission.
Because detectable HI is soextensive, HI is an exceptionally sensitive tracer of tidalinteractions between galaxies. Longstreamers and tails of HI trace the interaction histories ofpairs and groups of galaxies. See theHIRogues Gallery for somespectacular examples.
Optical (white)and HI (blue) imagesof the strongly interacting galaxies NGC 4038 and NGC 4039 (also knownas the 'antennae'). The velocity distributions of the long HItidal tails provide strong constraints for computer models of theinteraction history. Imagecredit
Theradio continuum (red) and HI(blue) images of the post-merger pair of galaxies UGC 813 and UGC 816indicate that the disks of these two galaxies passed through each otherabout 50 million years ago.Imagecredit
Another application of the HI spectraof galaxies isdetermining departures from smooth Hubble expansion in the localuniverse via the Tully-Fisherrelation.Most galaxies obeytheempirical luminosity-velocity relation (Tully, R. B., & Fisher, J.R. 1977, A&A, 54, 661):
$$L propto v_{rm m}^4~,$$ where $v_{rm m}$ is the maximumrotation speed. Arguments based onthe virial theorem can explain the Tully-Fisher relation if allgalaxies havesame central mass density and density profile, differing only in scalelength, and also have the same mass-to-light ratio. Thus a measurementof $v_{rm m}$ yields an estimate of $L$ that is independent of theHubble distance $D_{rm H}$. The Tully-Fisher distance $D_{rm TF}$ canbe calculated from this 'standard candle' $L$ and the apparentluminosity. Apparentluminosities in the near infrared ($lambda sim 2,mu$m) are favoredbecause the near-infrared mass-to-light ratio of stars is nearlyconstant and independent of the star-formation history, and becauseextinction by dust is much less than at optical wavelengths.Differences between $D_{rm TF}$ and $D_{rm H}$ are ascribed to the peculiarvelocities of galaxiescaused by intergalactic gravitationalinteractions. The magnitudes and scale lengths of the peculiar velocitydistributions are indications of the average density and clumpiness ofmass on megaparsec scales.
Large-Scale Structure
The Great Attractor.
BaryonAcoustic Oscillations
The WMAP five-year all-skythe CMB.Imagecredit
The WMAP five-year power spectrum ofthe CMB.Imagecredit
Epochof Reionization (EOR)
Most of the normal matter in the early universe was fully ionizedhydrogen and helium gas, plus trace amounts of heavier elements. This smoothly distributed gas cooled as the universe expanded, and thefree protons and electrons recombined to form neutral hydrogen at aredshift $z approx 1091$ when the age of the universe was about $3.8times 10^5$ years. The hydrogen remained neutral during the'dark ages' prior to the formation of the first ionizing astronomicalsources, massive ($M > 100 M_odot$) stars, galaxies, quasars, andclusters of galaxies, by gravitational collapse ofoverdense regions. These astronomical sources gradually startedreionizing the universe when it was several hundred million years old($z sim 10$) and completely reionized the universe by the time it wasabout $10^9$ years old ($z sim 6$). This era is called the epochof reionization.
Thehighly redshift HI signal was uniform prior to the firstreionization. The signal developed structure on angular scales upto several arcmin as the first sources created bubbles of ionizedhydrogen around them. As the bubbles grew and merged, the HIsignal developed frequency structure corresponding to the redshifted HIline frequency. The characteristic size of the larger bubblesreached about 10 Mpc at $z sim 6$, corrresponding to HI signals havingangular scales of several arcmin and covering frequency ranges ofseveral MHz. These the HI signals encode unique informationabout the formation of the earliest astronomical sources.
TheHI signals produced by the EOR will be very difficult to detect becausethey are weak (tens of mK), relatively broad in frequency, redshiftedto low frequencies (< 200MHz) plagued by radio-frequency interference, and lie behind a muchbrighter(tens of K) foreground of extragalactic continuum radio sources. Nonetheless,the potential scientific payoff is so great that several groups aroundthe world are developing instruments to detect the HI signature of theEOR. One such instrument is thePrecision Array to Probe the Epoch of Reionization (PAPER), a jointproject of UC Berkeley, the NRAO, and the University of Virginia thathas placed a test array in Green Bank, WV and will soon construct asciencearray in western Australia, where there is very little RFI.
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