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astro-ph0002001
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QSO Absorption Line Constraints on Intragroup High--Velocity Clouds
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[
{
"author": "Jane~C.~Charlton\\altaffilmark{1}"
},
{
"author": "Christopher~W.~Churchill"
},
{
"author": "and Jane~R.~Rigby"
}
] |
We show that the number statistics of moderate redshift {\MgII} and Lyman limit absorbers may rule out the hypothesis that high velocity clouds are infalling intragroup material.
|
[
{
"name": "astro-ph0002001.tex",
"string": "\\documentstyle[emulateapj,rotate,times]{article}\n%\\documentstyle[11pt,rotate,aaspp4]{article}\n%\\documentstyle[12pt,aasms4]{article}\n \n% this is the modification \n \n%-----------------------------------------------------------------------------\n%\n% personal LaTex mark--ups\n%\n%\n% **** DO NOT MODIFY BELOW **** **** DO NOT MODIFY BELOW **** \n%\n\\newcommand{\\etal}{et~al.\\ }\n\\newcommand{\\PVdblt}{{\\rm P}\\kern 0.1em{\\sc v}~$\\lambda\\lambda 1117, 1128$}\n\\newcommand{\\CaIIdblt}{{\\rm Ca}\\kern 0.1em{\\sc ii}~$\\lambda\\lambda 3934, 3969$}\n\\newcommand{\\AlIIIdblt}{{\\rm Al}\\kern 0.1em{\\sc iii}~$\\lambda\\lambda 1854, 1862$}\n\\newcommand{\\CIVdblt}{{\\rm C}\\kern 0.1em{\\sc iv}~$\\lambda\\lambda 1548, 1550$}\n\\newcommand{\\MgIIdblt}{{\\rm Mg}\\kern 0.1em{\\sc ii}~$\\lambda\\lambda 2796, 2803$}\n\\newcommand{\\NVdblt}{{\\rm N}\\kern 0.1em{\\sc v}~$\\lambda\\lambda 1238, 1242$} \n\\newcommand{\\SVIdblt}{{\\rm S}\\kern 0.1em{\\sc vi}~$\\lambda\\lambda 933, 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Fe}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\FeIII}{\\hbox{{\\rm Fe}\\kern 0.1em{\\sc iii}}}\n\\newcommand{\\MnII}{\\hbox{{\\rm Mn}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\MgI}{\\hbox{{\\rm Mg}\\kern 0.1em{\\sc i}}}\n\\newcommand{\\MgII}{\\hbox{{\\rm Mg}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\MgIII}{\\hbox{{\\rm Mg}\\kern 0.1em{\\sc iii}}}\n\\newcommand{\\NI}{\\hbox{{\\rm N}\\kern 0.1em{\\sc i}}}\n\\newcommand{\\NII}{\\hbox{{\\rm N}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\NIII}{\\hbox{{\\rm N}\\kern 0.1em{\\sc iii}}}\n\\newcommand{\\NV}{\\hbox{{\\rm N}\\kern 0.1em{\\sc v}}}\n\\newcommand{\\OVI}{\\hbox{{\\rm O}\\kern 0.1em{\\sc vi}}}\n\\newcommand{\\OI}{\\hbox{{\\rm O}\\kern 0.1em{\\sc i}}}\n\\newcommand{\\OII}{\\hbox{[{\\rm O}\\kern 0.1em{\\sc ii}]}}\n\\newcommand{\\OIV}{\\hbox{{\\rm O}\\kern 0.1em{\\sc iv}]}}\n\\newcommand{\\SI}{{\\rm S}\\kern 0.1em{\\sc i}}\n\\newcommand{\\SIV}{{\\rm S}\\kern 0.1em{\\sc iv}}\n\\newcommand{\\SVI}{{\\rm S}\\kern 0.1em{\\sc vi}}\n\\newcommand{\\SiI}{\\hbox{{\\rm Si}\\kern 0.1em{\\sc i}}}\n\\newcommand{\\SiII}{\\hbox{{\\rm Si}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\SiIII}{\\hbox{{\\rm Si}\\kern 0.1em{\\sc iii}}}\n\\newcommand{\\SiIV}{\\hbox{{\\rm Si}\\kern 0.1em{\\sc iv}}}\n\\newcommand{\\SII}{\\hbox{{\\rm S}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\SIII}{\\hbox{{\\rm S}\\kern 0.1em{\\sc iii}}}\n\\newcommand{\\NaI}{\\hbox{{\\rm Na}\\kern 0.1em{\\sc i}}}\n\\newcommand{\\TiII}{\\hbox{{\\rm Ti}\\kern 0.1em{\\sc ii}}}\n\\newcommand{\\kms}{\\hbox{km~s$^{-1}$}}\n\\newcommand{\\cmsq}{\\hbox{cm$^{-2}$}}\n\\newcommand{\\cc}{\\hbox{cm$^{-3}$}}\n \n%\n% **** DO NOT MODIFY ABOVE **** **** DO NOT MODIFY ABOVE **** \n%\n%-----------------------------------------------------------------------------\n%\\clubpenalty=450\n%\\widowpenalty=1000000\n\\tighten\n\\begin{document}\n \n%\\received{date month year}\n%\\accepted{date month year}\n%\\journalid{number}{date month year}\n%\\articleid{number}{number}\n%\\slugcomment{submitted to: {\\it The Astrophysical Journal Letters}}\n \n\\lefthead{CHARLTON ET~AL.}\n\\righthead{CONSTRAINTS ON INTRAGROUP HVCs}\n\n\n%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%\n\n\\title{QSO Absorption Line Constraints on Intragroup High--Velocity\nClouds}\n\n\\author{Jane~C.~Charlton\\altaffilmark{1}, Christopher~W.~Churchill,\nand Jane~R.~Rigby}\n\\affil{The Pennsylvania State University, University Park, PA 16802 \\\\\ncharlton, cwc, [email protected]}\n\n\\altaffiltext{1}{Center for Gravitational Physics and Geometry}\n\n\\begin{abstract}\nWe show that the number statistics of moderate redshift {\\MgII} and\nLyman limit absorbers may rule out the hypothesis that high velocity\nclouds are infalling intragroup material.\n\\end{abstract}\n\n\\keywords{quasars: absorption lines --- intergalactic\nmedium --- Local Group}\n\n\n%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%\n\n\\section{Introduction}\n\nThe origin(s) of high--velocity clouds (HVCs), gaseous material that\ndeparts from the Galactic rotation law by more than 100~{\\kms}, is a\ntopic under debate.\nUndoubtedly, some HVCs arise from tidal streams (e.g.\\ the Magellanic\nStream), and from fountain processes local to the Galaxy\n(\\cite{wakker97}). \nRecently, however, the hypothesis that {\\it most\\/} HVCs are\ndistributed ubiquitously throughout the Local Group and are relics of\ngroup formation has returned to favor (\\cite{blitz};\n\\cite{braun-burton}).\n\nIn the intragroup HVC hypothesis\n(1) the cloud kinematics follow the Local Group standard of\nrest (LGSR), not the Galactic standard of rest (GSR), with\nthe exception of some HVCs related to tidal stripping or\nGalactic fountains (\\cite{blitz}; \\cite{braun-burton}); \n(2) the cloud Galactocentric distances are typically 1~Mpc; \n(3) the extended HVC cloud complexes are presently\naccreting onto the Milky Way;\n(4) the clouds are local analogs of the Lyman limit absorbers observed in\nquasar spectra; \n(5) the clouds have masses of $10^{7}$~M$_{\\odot}$ and greater; and \n(6) the metallicities are lower than expected if the material\noriginated from blowout or fountains from the Milky Way\n(\\cite{wakker99}; \\cite{bowen93}).\n\nBlitz \\etal (1999\\nocite{blitz}; hereafter BSTHB) suggest that there\nare $\\sim 300$ clouds above the $21$--cm $N({\\HI})$ detection threshold of \n$\\simeq 2 \\times 10^{18}$~{\\cmsq}. These clouds have radii $\\sim 15$~kpc \nand are ubiquitous throughout the group. \n\nBraun and Burton (1999\\nocite{braun-burton}, hereafter BB) \ncataloged $65$ Local Group CHVCs, which represent a \nhomogeneous subset of the HVC population discussed by BSTHB.\nHigh resolution $21$--cm observations (\\cite{bb-hires})\nshow that the CHVCs, have compact, $N({\\HI})>10^{19}$~{\\cmsq},\ncold cores with few {\\kms} FWHM surrounded by extended \n``halos'' with FWHM $\\sim 25$~{\\kms}.\nThe typical radius is $\\sim5$--$8$~kpc at the estimated distance\nof $\\sim700$~kpc. The BB sample is homogeneous but is not\ncomplete; they estimate that there could be as many as\n$200$ Local Group CHVCs.\n\nRecently, Zwaan and Briggs (2000\\nocite{zwaan}) reported\nevidence in contradiction of the intragroup hypothesis.\nIn a blind {\\HI} $21$--cm survey of extragalactic groups, sensitive to\n$N({\\HI}) \\sim 10^{18}$~{\\cmsq} (capable of detecting $\\simeq\n10^{7}$~M$_{\\odot}$ {\\HI} clouds), they failed to locate any\nextragalactic counterparts of the Local Group HVCs.\nThis is in remarkable contrast to the numbers predicted.\nIf intragroup HVCs exist around all galaxies or galaxy groups, and\nthe {\\HI} mass function is the same in extragalactic groups as measured\nlocally, then Zwaan \\& Briggs should have detected $\\sim70$ in groups\nand $\\sim 250$ around galaxies ($\\sim 10$ and $\\sim 40$ for the CHVCs,\nrespectively).\n\nThus, the Zwaan and Briggs\\nocite{zwaan} result is in conflict with\nthe intragroup HVC hypothesis.\nSince the hypothesized intragroup clouds are remnants of galaxy\nformation and are shown to be stable against destruction mechanisms\n(BSTHB\\nocite{blitz}), they are predicted to form at very high redshifts \nand to be ubiquitous in galaxy groups to the present epoch.\nIn this {\\it Letter}, we argue that the version of the intragroup HVC \nhypothesis presented by BSTHB is also in conflict with the observed\nredshift number density of moderate redshift $(z \\simeq 0.5)$ {\\MgII} \nand Lyman limit (LLS) quasar absorption line systems.\nWe also find that the properties of the extragalactic analogs\nof the BB CHVCs are severely constrained.\n\nIn general, the redshift number density of a non--evolving population\nof objects, to be interpreted as the number per unit redshift, is written\n\\begin{equation}\n\\frac{dN}{dz} = C_{f}\\frac{n\\sigma c}{H_0} \n\\left( 1 + z \\right) \\left( 1 + 2q_{0}z \\right) ^{-1/2} , \n\\label{eq:dndz}\n\\end{equation}\nwhere $C_{f}$ is the covering factor, $n$ is the number density of\nabsorbing structures, and $\\sigma$ and $C_{f}$ are the surface area\npresented by each structure and its covering factor for detectable\nabsorption. Throughout, we use $H_{0} = 100$~{\\kms}~Mpc$^{-1}$ and $q_0\n= 0.5$, which gives $dN/dz \\propto (1+z)^{1/2}$.\n\n\\section{M\\lowercase{g} II Systems}\n\\label{sec:mgii}\n\nThe statistics of {\\MgII} absorbers are well--established at $0.3 \\leq\nz \\leq 2.2$. \nFor rest--frame equivalent widths of $W({\\MgII}) > 0.3$~{\\AA}\n(``strong'' {\\MgII} absorption)\nSteidel \\& Sargent (1992\\nocite{ss92}) found $dN/dz = 0.8\\pm 0.2$ for\n$z\\simeq 0.5$ with a redshift dependence consistent with no evolution\nexpectations.\nNormal, bright ($L \\geq 0.1~L^{\\ast}$) galaxies are almost\nalways found within $40$~kpc of strong {\\MgII} absorbers\n(\\cite{bb91}; \\cite{bergeron92}; \\cite{lebrun93}; \\cite{sdp94};\n\\cite{s95}; \\cite{3c336}).\nFrom the Steidel, Dickinson, and Persson survey, all but $3$ of $58$ strong {MgII} absorbers, \ndetected toward $51$ quasars, have identified galaxies with a coincident\nredshift within that impact parameter (sky projected separation\nfrom the quasar line of sight) (see \\cite{cc96}).\nAlso, it is rare to observe a galaxy with an impact parameter less than \n$\\sim 40 h^{-1}$~kpc that does {\\it not\\/} give rise to {\\MgII} \nabsorption with $W({\\MgII}) > 0.3$~{\\AA} (\\cite{s95}). \nIn $25$ ``control fields'' of quasars, without observed strong {\\MgII} \nabsorption in their spectra, only two galaxies had impact parameters\nless than $40 h^{-1}$~kpc (see also \\cite{cc96}).\nAs such, the regions within $\\sim 40 h^{-1}$~kpc of typical galaxies\naccount for the vast majority of {\\MgII} absorbers above this\nequivalent width threshold; there is nearly a ``one--to--one''\ncorrespondence.\nIf we accept these results, it would imply that there is little\nroom for a contribution to $dN/dz$ from a population of intragroup\nclouds {\\it which have impact parameters much greater than $\\sim40 h^{-1}$~kpc}.\n\nHowever, the predicted cross section for {\\MgII} absorption from \nthe extragalactic intragroup clouds analogous to HVCs would be substantial.\nWe quantify the overprediction of the redshift path density by\ncomputing the ratio of $dN/dz$ of the intragroup clouds to that of\n{\\MgII} absorbing galaxies,\n\\begin{equation}\n\\frac{(dN/dz)_{cl}}{(dN/dz)_{gal}} = F \n \\left( \\frac{f_{cl}}{f_{gal}} \\right)\n \\left( \\frac{N_{cl}}{N_{gal}} \\right) \n \\left( \\frac{R_{cl}}{R_{gal}} \\right) ^{2} ,\n\\label{eq:ratio}\n\\end{equation}\nwhere $F$ is the fraction of {\\MgII} absorbing galaxies that reside in\ngroups having intragroup HVC--like clouds, $f_{cl}$ is the fraction\nof the area of the clouds and $f_{gal}$ is the fraction of the area\nof the galaxies that would produce $W({\\MgII}) > 0.3$~{\\AA} along the \nline of sight, and $N_{cl}$ and $N_{gal}$ are the number of clouds \nand galaxies per group, respectively.\nThe cross section of the group times the intragroup cloud covering\nfactor, $C_{f}\\cdot \\pi R^{2}_{gr}$, is equal to $N_{cl} \\cdot \\pi R^{2}_{cl}$.\nThe total predicted $dN/dz$ for {\\MgII} absorbers with $W({\\MgII}) >\n0.3$~{\\AA} is then,\n\\begin{equation}\n\\left( \\frac{dN}{dz} \\right) _{tot} =\n\\left( \\frac{dN}{dz} \\right) _{gal}\n\\left[ 1 + \\frac{(dN/dz)_{cl}}{(dN/dz)_{gal}} \\right] .\n\\label{eq:totaldndz}\n\\end{equation}\nIf virtually all {\\MgII} absorbers are accounted for by galaxies,\nit is required that $(dN/dz)_{tot} \\simeq (dN/dz)_{gal}$; the left\nhand side of Equation~\\ref{eq:ratio} must be very close to zero.\n\nIn the BSTHB version of the intragroup HVC model, the ``best''\nexpected values are $N_{cl}=300$ and $R_{cl} = 15$~kpc\n(BSTHB; \\cite{blitz-privcomm}); \nif we take $R_{gal} = 40$~kpc and $f_{gal} = 1$ (\\cite{s95}),\nand assuming $N_{gal} = 4$, we find that the covering factor for\n{\\MgII} absorption from extragalactic analogs to the Local Group HVCs\nwould exceed that from galaxies by a factor of $\\sim 10$ for $F=1$ and\n$f_{cl}=1$, giving $(dN/dz)_{tot} \\simeq 9$.\nMore recently, Blitz and Robinshaw (2000\\nocite{blitzdsph}) have\nsuggested that sizes may be smaller ($R_{cl} = 8$~kpc) when \nbeam--smearing is considered.\nConsidering this as an indication of uncertainties in the\nBSTHB parameters, and considering $2 < N_{gal} < 6$ for the\ntypical number of group galaxies, we find, for $F=1$ and\n$f_{cl}=1$, a range $2 < (dN/dz)_{cl}/(dN/dz)_{gal} < 21$.\nThis corresponds to $2.6 < (dN/dz)_{tot} < 17.6$.\nIt is unlikely that $F$ is significantly less than unity;\nthe majority of galaxies reside in groups like the Local Group\nthat would have HVC analogs.\nIn order that $(dN/dz)_{tot} \\sim (dN/dz)_{gal}$,\n$f_{cl} \\ll 0.2$ is required.\n\nIt is not clear what fraction $f_{cl}$ of\nHVCs with $N({\\HI})$ above the $10^{18}$~{\\cmsq} detection \nthreshold will give rise to $W({\\MgII}) \\geq 0.3$~{\\AA}\nbecause the equivalent width is sensitive\nto the metallicity and internal velocity dispersion of the clouds.\nBased upon Cloudy (\\cite{ferland}) photoionization equilibrium models, \na cloud with $N({\\HI}) \\simeq 10^{18}$~{\\cmsq}, subject to\nthe ionizing metagalactic background (\\cite{haardt-madau}), would give\nrise to {\\MgII} absorption with $N({\\MgII}) \\simeq\n10^{14}N_{18}(Z/Z_{\\odot})$~{\\cmsq}, where $N_{18}$ is the {\\HI}\ncolumn density in units of $10^{18}$~{\\cmsq} and $Z/Z_{\\odot}$ is the\nmetallicity in solar units.\nFor optically thick clouds, those with $N({\\HI}) \\geq\n10^{17.5}$~{\\cmsq}, this result is insensitive to the assumed\nionization parameter\\footnote{The ionization parameter is the ratio of\nthe number density of hydrogen ionizing photons to the number density\nof electrons, $n_{\\gamma}/n_{e}$.} over the range $10^{-4.5}$ to\n$10^{-1.5}$.\n\nBSTHB expect HVCs to have $Z/Z_{\\odot} \\sim 0.1$.\nFor $N_{18} = 2$ and $Z/Z_{\\odot} = 0.1$, clouds with internal\nvelocity dispersions of $\\sigma _{cl} \\geq 20$~{\\kms} (Doppler\n$b \\geq 28$~{\\kms}) give rise to $W({\\MgII}) \\geq 0.5$~{\\AA}.\nFor $\\sigma_{cl} = 10$~{\\kms} ($b = 14$~{\\kms}), $W({\\MgII}) = 0.3$~{\\AA}.\nThe CHVC ``halos'' typically have FWHM of $29$--$34$~{\\kms}, which\ncorresponds to $\\sigma _{cl} \\sim 12$--$14$~{\\kms} \n(\\cite{braun-burton}). \nThus it appears that most lines of sight through the BSTHB\nextragalactic analogs will produce strong {\\MgII} absorption.\nCertainly $f_{cl} > 0.2$, so there is a serious\ndiscrepancy between the predicted $(dN/dz)_{tot}$ and the\nobserved value.\n\nHowever, if the intragroup clouds have lower metallicities, this would\nresult in smaller $W({\\MgII})$. Unfortunately, there has only been\none metallicity estimate published for an HVC, which may or may not be\nrelated to the Galaxy. Braun and Burton (2000\\nocite{bb-hires}) estimate that\nCHVC 125+41-207, with $W({\\MgII}) = 0.15$~{\\AA}, has a metallicity of $0.04\n< Z/Z_{\\odot} < 0.07$, however this is quite uncertain because of the \neffects of beam smearing on measuring the $N({\\HI})$ value.\nBecause of the uncertainties, we simply state that a population of low\nmetallicity clouds could reduce the discrepancy\nbetween the predicted redshift density for intragroup clouds,\n$(dN/dz)_{cl}$, and the observed value of $(dN/dz)_{tot}$. However,\nthen the expected number of smaller $W({\\MgII})$ systems to arise from\nintragroup clouds would be increased.\n\nThe observed {\\MgII} equivalent width distribution rises rapidly below\n$0.3$~{\\AA} (``weak'' {\\MgII} absorbers), such that $dN/dz =\n2.2\\pm0.3$ for $W({\\MgII}) > 0.02$~{\\AA} at $z=0.5$ (\\cite{weak1}).\nTo this equivalent width limit, {\\MgII} absorption could be observed\nfrom intragroup HVCs with $N_{18}=2$ and metallicities as low as\n$Z/Z_{\\odot} = 0.0025$ [for $N({\\MgII}) = 10^{11.7}$~{\\cmsq},\n$W({\\MgII})$ is independent of $\\sigma _{cl}$]. However, almost all\n($9$ out of a sample of $10$) {\\MgII} absorbers with $W({\\MgII}) <\n0.3$~{\\AA} do {\\it not\\/} have associated Lyman limit breaks\n(\\cite{paper1}); that is, their $N({\\HI})$ is more than a decade below\nthe sensitivity of $21$--cm surveys. Thus, based upon available data,\nroughly $90$\\% of the ``weak'' {\\MgII} absorbers do not have the\nproperties of HVCs, and therefore are {\\it not\\/} analogous to the\nclouds invoked for the intragroup HVC scenario. If $10$\\% of the weak\n{\\MgII} absorbers are analogs to the intragroup HVCs, they would\ncontribute an additional $0.20$ to $(dN/dz)_{cl}$.\n\nSince the BB CHVC extragalactic analogs have smaller cross sections,\nwe should separately consider whether they would produce a discrepancy\nwith the observed {\\MgII} absorption statistics. BB observed\n$N_{cl}=65$ and inferred a typical $R_{cl} = 5$--$8$~kpc for the\nCHVCs, however a complete sample might have $N_{cl}=200$. Assuming\n$N_{gal}=2$--$6$, $R_{gal} = 40$~kpc, $f_{gal} = 1$, and $F=1$, for\nthe BB subsample of the HVC population, we obtain \n$0.17 f_{cl} < (dN/dz)_{cl}/(dN/dz)_{gal} < 4.0 f_{cl}$.\n\nThe cores of the CHVCs have $N({\\HI}) > 10^{19}$~{\\cmsq} and they\noccupy only about $15$\\% of the detected extent. \nFor $Z/Z_{\\odot} > 0.01$ and $\\sigma_{cl} = 10$~{\\kms}, these cores can produce\n$W({\\MgII}) \\ge 0.3$~{\\AA} over their full area. It follows that\n$f_{cl} = 0.15$, which yields $0.025 < (dN/dz)_{cl}/(dN/dz)_{gal} <\n0.6$. Depending on the specific parameters, there may or may not be a\nconflict with the observed $(dN/dz)_{tot}$ for strong {\\MgII}\nabsorption.\n\nThe ``halos'' of the CHVCs have $N({\\HI}) > 10^{18}$~{\\cmsq} and, as\ndiscussed above, would produce weak {\\MgII} absorption for\n$Z/Z_{\\odot} > 0.005$ over most of the cloud area. This implies\ncontribution to $(dN/dz)$ from BB CHVC analogs that is in the range\n$0.14 < (dN/dz)_{cl} < 3.2$. If the number were at the high end of\nthis range, the cross section would be comparable to the observed\n$(dN/dz)$ for weak {\\MgII} absorbers at $z=0.5$. However, as noted\nabove when considering the BSTHB scenario, there is a serious\ndiscrepancy. Only $\\sim 10$\\% of the weak {\\MgII} absorbers show a\nLyman limit break, so extragalactic analogs of the BB CHVCs can only\nbe a fraction of the weak {\\MgII} population. Regions of CHVCs at larger\nradii, with $N({\\HI})$ below the threshold of present $21$--cm\nobservations, are constrained to have $Z/Z_{\\odot} \\ll 0.01$ in order\nthat they do not produce a much larger population of weak {\\MgII}\nabsorbers with Lyman limit break than is observed.\n\n\\section{Lyman Limit Systems}\n\\label{sec:lls}\n\nThe redshift number density of LLS also places strong constraints on\nintragroup environments that give rise to Lyman breaks in quasar\nspectra. This argument is not sensitive to the assumed\ncloud velocity dispersion and/or metallicity.\n\nStatistically, $dN/dz$ for {\\MgII} systems is consistent\n(1~$\\sigma$) with $dN/dz$ for LLS.\nAt $z \\simeq 0.5$, LLS have $dN/dz = 0.5 \\pm 0.3$ (\\cite{kplls}) and\n{\\MgII} systems have $dN/dz = 0.8 \\pm 0.2$ (\\cite{ss92}).\nChurchill \\etal (2000a\\nocite{paper1}) found a Lyman limit break\n[i.e.\\ $N({\\HI}) \\geq 10^{16.8}$~{\\cmsq}] for each system in a sample\nof ten having $W({\\MgII}) > 0.3$~{\\AA}.\nLLS and {\\MgII} absorbers have roughly the same redshift number\ndensity and therefore {\\MgII}--LLS absorption must almost always arise \nwithin $\\sim 40$~kpc of galaxies (\\cite{s93}).\nAs such, there is little latitude for a substantial contribution\nto $dN/dz$ from intragroup Lyman limit clouds.\n\nUsing Equation~\\ref{eq:dndz}, we could estimate this contribution by\nconsidering the volume density of galaxy groups and the cross section\nfor HVC Lyman limit absorption in each.\nHowever, the volume density of groups is not well measured,\nparticularly out to $z=0.5$.\n\nInstead, we make a restrictive argument based upon a comparison\nbetween the cross sections for Lyman limit absorption of $L^{\\ast}$ \ngalaxies and for HVCs in a typical group (similar to the discussion of\n{\\MgII} absorbers in \\S~\\ref{sec:mgii}).\nAgain, we simply compare the values of $C_f$ for the different\npopulations of objects in a typical group.\nThe covering factor for HVCs within the group is\n\\begin{equation}\nC_f = N_{cl}\n\\frac{\\left(R_{cl}\\right)^2}{\\left(R_{gr}\\right)^2} .\n\\end{equation}\nThe best estimate for the BSTHB version of the intragroup HVC model,\nwith $N_{cl} = 300$, $R_{cl} = 15$~kpc, and a group radius \n$R_{gr} = 1.5$~Mpc, gives $C_{f} = 0.03$ for $N({\\HI}) \\simeq 2 \\times\n10^{18}$~{\\cmsq}. \nIf instead we use the BB number of observed CHVCs, $N_{cl} = 65$,\nand $R_{cl} = 5$--$8$~kpc, we obtain a much smaller number, \n$0.0007 < C_{f} < 0.0018$.\nHowever, if the BB sample is corrected for incompleteness such that\n$N_{cl} = 200$, these numbers increase so that\n$0.002 < C_{f} < 0.006$.\n\nIn comparison, a typical group with $\\sim4$ $L^{\\ast}$ galaxies, each with a\nLyman limit absorption cross section of $R_{cl} \\sim 40$~kpc, would\nhave $C_f = 0.002$. \nIf they existed with the properties discussed,\nthe extragalactic analogs to the BSTHB HVCs would dominate the\ncontribution of $L^{\\ast}$ galaxies to the $dN/dz$ of LLS by at least\na factor of $\\sim 15$, and this is only considering HVC regions with\n$N({\\HI}) > 2 \\times 10^{18}$~{\\cmsq} that are detected in the $21$--cm\nsurveys. \nAny extensions in area below this threshold value [down to\n$N({\\HI}) \\sim 5 \\times 10^{16}$~{\\cmsq}] would worsen the discrepancy.\nAs such, the BSTHB hypothesis is definitively ruled out.\n\nFor regions of BB CHVCs with $N({\\HI}) > 2 \\times 10^{18}$~{\\cmsq}, the\ncovering factor ranges from $C_f =0.0007$ to $C_f=0.006$ depending\non assumed sizes and corrections for incompleteness. This ranges\nfrom $35$--$300$\\% of the cross section for the $L^{\\ast}$ galaxies.\nThe {\\it total} observed $dN/dz$ for Lyman limit absorption \n(down to $\\log N({\\HI}) = 17$~{\\cmsq}) is only $\\sim 0.5$, even a $35$\\% \ncontribution to the Lyman limit cross section from HVCs \nthat are separate from galaxies creates a discrepancy.\nThis would imply that the result that most lines of sight within \n$40$~kpc of a typical $L^{\\ast}$ galaxy produce Lyman limit\nabsorption is incorrect.\nThis would further imply that there is a substantial population of\nstrong {\\MgII} absorbers without Lyman limit breaks\n(to account for $dN/dz = 0.8$ for strong {\\MgII}\nabsorption) or of strong {\\MgII} absorbers not associated\nwith galaxies. \nBoth types of objects are rarely observed\n(\\cite{paper1}; \\cite{bb91}; \\cite{bergeron92}; \\cite{lebrun93};\n\\cite{sdp94}; \\cite{s95}; \\cite{3c336}).\nFurthermore, $C_f=0.002$ for BB CHVCs only takes into account the\nfraction of the BB CHVC areas with $N({\\HI}) > 2 \\times 10^{18}$~{\\cmsq}.\nTherefore, the extended ``halos'' around the CHVCs are also\nconstrained not to contribute substantial cross section for Lyman\nlimit absorption along extragalactic lines of sight.\n\n\\section{Summary}\n\\label{sec:summary}\n\nWe have made straight--forward estimates of the\npredicted redshift number density at $z \\simeq 0.5$ of\n{\\MgII} and LLS absorption from hypothetical extragalactic\nanalogs to intragroup HVCs as expected by extrapolating from the BSTHB\nand BB Local Group samples. We find that it is difficult to reconcile\nthe intragroup hypothesis for HVCs with the observed $dN/dz$ of\n{\\MgII} and LLS systems.\n\nThe discrepancy between the $dN/dz$ of {\\MgII}--LLS absorbers and the\nobserved covering factor of ``intragroup'' HVCs could be reduced if\nthe HVCs have a clumpy structure. Such structure would result in\n{\\MgII}--LLS absorption observable only in some fraction, $f_{los}$,\nof the lines of sight through the cloud. Effectively, this reduces\nthe covering factor for Lyman limit absorption, or the value of\n$f_{cl}$ in equation (2) for {\\MgII} absorbers. Considering beam smearing\nin $21$--cm surveys, substructures would be detected above a\n$N({\\HI}) > 2 \\times 10^{18}$~{\\cmsq} $21$--cm detection threshold\nif their column densities were\n$N({\\HI})_{sub} > 2 \\times 10^{18}/f_{los}$~{\\cmsq}. The predicted\n$dN/dz$ for HVC--like clouds could be reduced by a factor of ten if\n$f_{los} \\leq 0.1$, giving $N({\\HI})_{sub} > 2 \\times\n10^{19}$~{\\cmsq}. All the gas outside these higher {\\HI} column\ndensity substructures would need to be below the Lyman limit or the\narguments in \\S~\\ref{sec:lls} would hold. It is difficult to\nreconcile such a density distribution with the high resolution\nobservations of BB CHVCs which show diffuse halos around the\ncore concentrations (\\cite{bb-hires}), but these ideas merit\nfurther consideration.\n\n\\subsection{The BSTHB Scenario}\n\nWe conclude that the predicted $dN/dz$ from the hypothetical\npopulation of intragroup HVCs along extragalactic sight lines to\nquasars from the BSTHB scenario would exceed: \n\n1) the $dN/dz$ of {\\MgII} absorbers with $W({\\MgII}) \\geq 0.3$~{\\AA}.\nThis class of absorber is already known to arise within $\\sim 40 h^{-1}$~kpc\nof normal, bright galaxies (\\cite{bb91}; \\cite{bergeron92};\n\\cite{lebrun93}; \\cite{sdp94}; \\cite{s95}; \\cite{3c336}).\n\n2) the $dN/dz$ of ``weak'' {\\MgII} absorbers with $0.02 < W({\\MgII}) <\n0.3$~{\\AA} absorption. \nIn principle, weak {\\MgII} absorption could\narise from low metallicity, $0.005 \\leq Z/Z_{\\odot} < 0.1$, intragroup\nHVCs. However, the majority of observed weak systems are already\nknown to be higher metallicity, $Z/Z_{\\odot} \\simeq 0.1$, sub--Lyman\nlimit systems (\\cite{weak1}; \\cite{rigby}).\n\n3) the $dN/dz$ of Lyman limit systems.\nThese would be produced by all extragalactic BSTHB HVC analogs\nregardless of metallicity.\nHowever, {\\it most\\/} Lyman limit systems are seen to arise within \n$\\simeq 40 h^{-1}$~kpc of luminous galaxies (\\cite{s93}; \\cite{paper2}).\n\nThese points do not preclude a population of infalling intragroup\nclouds which do not present a significant cross section for\n$21$--cm absorption, as predicted by CDM models\n(\\cite{iforget}; \\cite{moore}).\nIn fact, such intragroup objects could be related to sub--Lyman limit\nweak {\\MgII} absorbers (\\cite{rigby}).\n\n\\subsection{The BB Scenario}\n\nThe properties of the BB CHVC population are also significantly\nconstrained by {\\MgII} and Lyman limit absorber statistics:\n\n1) They {\\it could\\/} produce $W({\\MgII}) \\geq 0.3$~{\\AA} in excess of\nwhat is observed if a large incompleteness correction is applied\n(i.e. so that $N_{cl} = 200$), or if relatively large sizes ($R_{cl}\n\\sim 8$~kpc) are assumed.\n\n2) They would be expected to contribute to the $dN/dz$ of\nweak [$W({\\MgII}) > 0.02$~{\\AA}] {\\MgII} absorption.\nHowever, based upon observations (\\cite{paper1}), only $\\sim 10$\\% \nof the population of weak {\\MgII} absorbers have Lyman--limit breaks.\nTherefore, only a small fraction of weak {\\MgII} absorption could \narise in extragalactic BB CHVC analogs.\n\n3) The $dN/dz$ for Lyman limit absorption from the hypothesized BB\nCHVC population could be a significant\nfraction, or comparable to that expected from the local environments of\n$L^{\\ast}$ galaxies (within $40$~kpc); the observed value is already\nconsistent with that produced by the galaxies.\n\n4) The CHVCs are observed to have a cool core with $N({\\HI}) >\n10^{19}$~{\\cmsq}, surrounded by a halo which typically extends to\n$R_{cl} \\sim 5$~kpc. It is natural to expect that the {\\HI} extends\nout to larger radii at smaller $N({\\HI})$ and should produce a Lyman\nlimit break out to the radius at which $N({\\HI}) < 10^{16.8}$~{\\cmsq}.\nAlthough there is expected to be a sharp edge to the {\\HI} disk at\n$N({\\HI}) \\sim 10^{17.5}$ or $10^{18}$~{\\cmsq} (\\cite{maloney};\n\\cite{corbelli}; \\cite{dove}), physically we would expect that this\nedge would level off at $\\sim 10^{17.5}$~{\\cmsq}, such that a\nsignificant cross section would be presented at $10^{16.8} < N({\\HI})\n< 10^{17.5}$~{\\cmsq}. Another possibility is that there is a sharp\ncutoff of the structure at $N({\\HI}) \\sim 2 \\times 10^{18}$~{\\cmsq},\nbut this is contrived.\n\n\\section{Conclusion}\n\nWe are forced to the conclusion that there can only be a limited\nnumber of extragalactic infalling group HVC analogs at $z \\sim 0.5$.\nFuture data could force a reevaluation of the relationships between\ngalaxies, Lyman limit systems, and {\\MgII} absorbers, but it seems\nunlikely that the more serious inconsistencies we have identified\ncould be reconciled in this way. A clumpy distribution of {\\HI}\ncould be constructed that would reduce the discrepancy, but would\nrequire very diffuse material (below the Lyman limit) around dense \ncores. Evolution in the population of HVCs\nis another possibility. If the extragalactic background radiation\ndeclined from $z=0.5$ to the present, the clouds would have been more\nionized in the past, and therefore would have had smaller cross\nsections at a given $N({\\HI})$. However, this does not explain why\nZwaan and Briggs (2000\\nocite{zwaan}) do not see the $z=0$\nextragalactic analogs to the HVCS or CHVCs. Our results are entirely\nconsistent with theirs, and the implications are the same: the\ndiscrepancies between the Local Group HVC population and the\nstatistics of {\\MgII} and Lyman limit absorbers can only be reconciled\nif most of the extragalactic HVC analogs are within $100$--$200$~kpc\nof galaxies, and not at large throughout the groups.\n\n\\acknowledgements\nWe thank L. Blitz, J. Bregman, J. Mulchaey, B. Savage, K. Sembach,\nT. Tripp, and especially B. Wakker and Buell Jannuzi, and our\nreferees for stimulating discussions and comments.\nSupport for this work was provided by NSF grant AST--9617185\n(J. R. R. was supported by an REU supplement) and by NASA grant NAG\n5--6399. \n\n%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%\n\n\n\\begin{thebibliography}{XXX}\n\n\\bibitem[Bergeron \\& Boiss\\'e 1991]{bb91}\nBergeron, J., \\& Boiss\\'e, P. 1991, A\\&A, 243, 344\n\n\\bibitem[Bergeron \\etal 1992]{bergeron92}\nBergeron, J., Cristiani, S., \\& Shaver, P. A. 1992, A\\&A, 257, 417\n\n\\bibitem[Blitz \\etal 1999]{blitz}\nBlitz, L., Spergel, D. N., Teuben, P. J., Hartmann, D., \\& Burton, W. B. 1999,\nApJ, 514, 818\n\n\\bibitem[Blitz 2000, private communication]{blitz-privcomm}\nBlitz, L. 2000, private communication\n\n\\bibitem[Blitz \\& Robinshaw 2000]{blitzdsph}\nBlitz, L., \\& Robinshaw, T. 2000, ApJ, submitted\n\n\\bibitem[Bowen \\& Blades 1993]{bowen93}\nBowen, D. V., \\& Blades, J. C. 1993, ApJ, 403, L55\n\n\\bibitem[Braun \\& Burton 2000]{bb-hires}\nBraun, R., \\& Burton, W. B. 2000, A\\&A, in press\n\n\\bibitem[Braun \\& Burton 1999]{braun-burton}\nBraun, R., \\& Burton, W. B. 1999, A\\&A, 341, 437\n\n\\bibitem[Charlton \\& Churchill 1998]{mgii-profiles}\nCharlton, J. C., \\& Churchill, C. W. 1998, ApJ, 499, 181\n\n\\bibitem[Charlton \\& Churchill 1996]{cc96}\nCharlton, J. C., \\& Churchill, C. W. 1996, ApJ, 465, 631\n\n\\bibitem[Churchill \\etal 2000a]{paper1}\nChurchill, C. W., Mellon, R. R., Charlton, J. C., Jannuzi, B. T.,\nKirhakos, S., Steidel, C. C., \\& Schneider, D. P. 2000a, ApJ,\nin press\n\n\\bibitem[Churchill \\etal 2000b]{paper2}\nChurchill, C. W., Mellon, R. R., Charlton, J. C., Jannuzi, B. T.,\nKirhakos, S., Steidel, C. C., \\& Schneider, D. P. 2000b, ApJ,\nsubmitted\n\n\\bibitem[Churchill \\etal 1999]{weak1}\nChurchill, C. W., Rigby, J. R., Charlton, J. C., \\& Vogt, S. S. 1999,\nApJS, 120, 51\n\n\\bibitem[Corbelli \\& Salpeter 1994]{corbelli}\nCorbelli, E. and Salpeter, E. E. 1994, \\apj, 419, 104\n\n\\bibitem[Dove \\& Shull 1994]{dove}\nDove, J. B., and Shull, J. M. 1994, \\apj, 423, 196\n\n\\bibitem[Ferland 1996]{ferland}\nFerland, G. J. 1996, Hazy, University of Kentucky Internal Report\n\n\\bibitem[Haardt \\& Madau 1996]{haardt-madau}\nHaardt, F., \\& Madau, P. 1996, ApJ, 461, 20\n\n\\bibitem[Klypin \\etal 1999]{iforget}\nKlypin, A. A., Kravtsov, A. V., Valenzuela, O., \\& Prada, F. 1999,\nApJ, submitted\n\n\\bibitem[Le~Brun \\etal 1993]{lebrun93}\nLe~Brun, V., Bergeron, J., Boiss\\'e, P., \\& Christian, C. 1993,\nA\\&A, 279, 33\n\n\\bibitem[Maloney 1993]{maloney}\nMaloney, P. 1993, \\apj, 414, 57\n\n\\bibitem[Moore \\etal 1999]{moore}\nMoore, B., Ghigna, S., Governato, F., Lake, G., Quinn, T.,\nStadel, J., \\& Tozzi, P. 1999,\nApJ, 524, L19\n\n\\bibitem[Rigby \\etal 2000]{rigby}\nRigby, J. R., Charlton, J. C., Churchill, C. W. 2000, ApJ, in preparation\n\n\\bibitem[Steidel 1993]{s93}\nSteidel, C. C. 1993, in The Environment and Evolution of Galaxy, eds.\\\nJ. M. Shull \\& H. A. Thronson, Jr., (Dordrecht: Kluwer Academic), 263\n\n\\bibitem[Steidel 1995]{s95}\nSteidel, C. C. 1995, in QSO Absorption Lines, ed. G. Meylan (Garching:\nSpringer--Verlag), 139\n\n\\bibitem[Steidel, Dickinson, \\& Persson 1994]{sdp94}\nSteidel, C. C., Dickinson, M. \\& Persson, E. 1994, ApJ, 437, L75 \n\n\\bibitem[Steidel \\etal 1997]{3c336}\nSteidel. C. C., Dickinson, M., Meyer, D. M., Adelberger, K. L., \\&\nSembach, K. R. 1997, ApJ, 480, 568\n\n\n\\bibitem[Steidel \\& Sargent 1992]{ss92}\nSteidel, C. C., \\& Sargent, W. L. W. 1992, ApJS, 80, 1\n\n\\bibitem[Stengler--Larrea \\etal 1995]{kplls}\nStengler--Larrea, E. A. \\etal 1995, ApJ, 444, 64\n\n\\bibitem[Wakker \\etal 1999]{wakker99}\nWakker, B. P., Howk, J. C., Savage, B. D., Tufte, S. L., Renolds,\nR. J., van Woerden, H., Schwarz, U. J., Peletier, R. F., \\& Kalberla,\nP. M. W. 1999, Nature, 400, 388\n\n\\bibitem[Wakker \\& van Woerden 1997]{wakker97}\nWakker, B. P., and van Woerden, H. 1997, ARA\\&A, 35, 509\n\n\\bibitem[Zwaan \\& Briggs 2000]{zwaan}\nZwaan, M. A., \\& Briggs, F. H. 2000, ApJL, in press\n\n\\end{thebibliography}\n\n\\end{document}\n\n"
}
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[
{
"name": "astro-ph0002001.extracted_bib",
"string": "\\begin{thebibliography}{XXX}\n\n\\bibitem[Bergeron \\& Boiss\\'e 1991]{bb91}\nBergeron, J., \\& Boiss\\'e, P. 1991, A\\&A, 243, 344\n\n\\bibitem[Bergeron \\etal 1992]{bergeron92}\nBergeron, J., Cristiani, S., \\& Shaver, P. A. 1992, A\\&A, 257, 417\n\n\\bibitem[Blitz \\etal 1999]{blitz}\nBlitz, L., Spergel, D. N., Teuben, P. J., Hartmann, D., \\& Burton, W. B. 1999,\nApJ, 514, 818\n\n\\bibitem[Blitz 2000, private communication]{blitz-privcomm}\nBlitz, L. 2000, private communication\n\n\\bibitem[Blitz \\& Robinshaw 2000]{blitzdsph}\nBlitz, L., \\& Robinshaw, T. 2000, ApJ, submitted\n\n\\bibitem[Bowen \\& Blades 1993]{bowen93}\nBowen, D. V., \\& Blades, J. C. 1993, ApJ, 403, L55\n\n\\bibitem[Braun \\& Burton 2000]{bb-hires}\nBraun, R., \\& Burton, W. B. 2000, A\\&A, in press\n\n\\bibitem[Braun \\& Burton 1999]{braun-burton}\nBraun, R., \\& Burton, W. B. 1999, A\\&A, 341, 437\n\n\\bibitem[Charlton \\& Churchill 1998]{mgii-profiles}\nCharlton, J. C., \\& Churchill, C. W. 1998, ApJ, 499, 181\n\n\\bibitem[Charlton \\& Churchill 1996]{cc96}\nCharlton, J. C., \\& Churchill, C. W. 1996, ApJ, 465, 631\n\n\\bibitem[Churchill \\etal 2000a]{paper1}\nChurchill, C. W., Mellon, R. R., Charlton, J. C., Jannuzi, B. T.,\nKirhakos, S., Steidel, C. C., \\& Schneider, D. P. 2000a, ApJ,\nin press\n\n\\bibitem[Churchill \\etal 2000b]{paper2}\nChurchill, C. W., Mellon, R. R., Charlton, J. C., Jannuzi, B. T.,\nKirhakos, S., Steidel, C. C., \\& Schneider, D. P. 2000b, ApJ,\nsubmitted\n\n\\bibitem[Churchill \\etal 1999]{weak1}\nChurchill, C. W., Rigby, J. R., Charlton, J. C., \\& Vogt, S. S. 1999,\nApJS, 120, 51\n\n\\bibitem[Corbelli \\& Salpeter 1994]{corbelli}\nCorbelli, E. and Salpeter, E. E. 1994, \\apj, 419, 104\n\n\\bibitem[Dove \\& Shull 1994]{dove}\nDove, J. B., and Shull, J. M. 1994, \\apj, 423, 196\n\n\\bibitem[Ferland 1996]{ferland}\nFerland, G. J. 1996, Hazy, University of Kentucky Internal Report\n\n\\bibitem[Haardt \\& Madau 1996]{haardt-madau}\nHaardt, F., \\& Madau, P. 1996, ApJ, 461, 20\n\n\\bibitem[Klypin \\etal 1999]{iforget}\nKlypin, A. A., Kravtsov, A. V., Valenzuela, O., \\& Prada, F. 1999,\nApJ, submitted\n\n\\bibitem[Le~Brun \\etal 1993]{lebrun93}\nLe~Brun, V., Bergeron, J., Boiss\\'e, P., \\& Christian, C. 1993,\nA\\&A, 279, 33\n\n\\bibitem[Maloney 1993]{maloney}\nMaloney, P. 1993, \\apj, 414, 57\n\n\\bibitem[Moore \\etal 1999]{moore}\nMoore, B., Ghigna, S., Governato, F., Lake, G., Quinn, T.,\nStadel, J., \\& Tozzi, P. 1999,\nApJ, 524, L19\n\n\\bibitem[Rigby \\etal 2000]{rigby}\nRigby, J. R., Charlton, J. C., Churchill, C. W. 2000, ApJ, in preparation\n\n\\bibitem[Steidel 1993]{s93}\nSteidel, C. C. 1993, in The Environment and Evolution of Galaxy, eds.\\\nJ. M. Shull \\& H. A. Thronson, Jr., (Dordrecht: Kluwer Academic), 263\n\n\\bibitem[Steidel 1995]{s95}\nSteidel, C. C. 1995, in QSO Absorption Lines, ed. G. Meylan (Garching:\nSpringer--Verlag), 139\n\n\\bibitem[Steidel, Dickinson, \\& Persson 1994]{sdp94}\nSteidel, C. C., Dickinson, M. \\& Persson, E. 1994, ApJ, 437, L75 \n\n\\bibitem[Steidel \\etal 1997]{3c336}\nSteidel. C. C., Dickinson, M., Meyer, D. M., Adelberger, K. L., \\&\nSembach, K. R. 1997, ApJ, 480, 568\n\n\n\\bibitem[Steidel \\& Sargent 1992]{ss92}\nSteidel, C. C., \\& Sargent, W. L. W. 1992, ApJS, 80, 1\n\n\\bibitem[Stengler--Larrea \\etal 1995]{kplls}\nStengler--Larrea, E. A. \\etal 1995, ApJ, 444, 64\n\n\\bibitem[Wakker \\etal 1999]{wakker99}\nWakker, B. P., Howk, J. C., Savage, B. D., Tufte, S. L., Renolds,\nR. J., van Woerden, H., Schwarz, U. J., Peletier, R. F., \\& Kalberla,\nP. M. W. 1999, Nature, 400, 388\n\n\\bibitem[Wakker \\& van Woerden 1997]{wakker97}\nWakker, B. P., and van Woerden, H. 1997, ARA\\&A, 35, 509\n\n\\bibitem[Zwaan \\& Briggs 2000]{zwaan}\nZwaan, M. A., \\& Briggs, F. H. 2000, ApJL, in press\n\n\\end{thebibliography}"
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astro-ph0002002
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Point source models for the gravitational lens B1608+656: Indeterminacy in the prediction of the Hubble constant
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[
{
"author": "Gabriela Surpi and Roger Blandford"
}
] |
We apply elliptical isothermal mass models to reproduce the point source properties, i.e. image positions, flux density ratios and time delay ratios, of the quadruple lens B1608+656. A wide set of suitable solutions is found, showing that models that only fit the properties of point sources are under-constrained and can lead to a large uncertainty in the prediction of H$_\circ$. We present two examples of models predicting H$_\circ\!=\!100{km\,s^{\!-\!1}Mpc^{\!-\!1}}$ ($\chi^2\!=\!4$) and H$_\circ\!=\!69{km\,s^{\!-\!1}Mpc^{\!-\!1}}$ ($\chi^2\!=\!24$).
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[
{
"name": "gsurpi.tex",
"string": "\\documentstyle[11pt,paspconf,epsf]{article}\n\n\\begin{document}\n\\title{Point source models for the gravitational lens B1608+656:\nIndeterminacy in the prediction of the Hubble constant}\n\n\\author{Gabriela Surpi and Roger Blandford}\n\\affil{California Institute of Technology 130-33, Pasadena CA 91125, USA}\n\n\\begin{abstract}\nWe apply elliptical isothermal mass\nmodels to reproduce the point source properties, i.e.\nimage positions, flux density ratios and\ntime delay ratios,\nof the quadruple lens B1608+656.\nA wide set of suitable solutions is found, showing that\nmodels that only fit the properties of point sources\nare under-constrained and can lead to a large\nuncertainty in the prediction of H$_\\circ$. We present\ntwo examples of models predicting H$_\\circ\\!=\\!100{\\rm\nkm\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$ ($\\chi^2\\!=\\!4$)\nand H$_\\circ\\!=\\!69{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$ ($\\chi^2\\!=\\!24$). \n\\end{abstract}\n\n\\keywords{gravitational lensing, models}\n\n\\vspace*{-0.4cm}\n\\section{Introduction}\n\nRelative positions, flux ratios and time delays of the 4\nimages in B1608+656 have been presented here by \nFassnacht (1999) and references therein ({\\it cf} Table~\\ref{obs}).\nKoopmans \\& Fassnacht (1999)\nhave concluded H$_\\circ\\!=\\!59^{+7}_{-6}{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$\nwithin the context of a family of parametrized, isothermal models. Here, we\ninvestigate whether a larger set of models allows a wider\nrange of Hubble constants.\n%\\vspace*{-0.1cm}\n\\section{Elliptical isothermal models}\n\nFollowing Blandford \\& Kundi\\'c (1997),\nwe adopt a scaled lensing potential $\\psi$ composed of two elliptical\ncontributions to describe the lensing galaxies G1 and G2 plus \nexternal shear $\\gamma$:\n\\vspace*{-0.2cm}\n\\begin{equation}\n\\psi_= \\sum_{i=1}^2 \\,\\, b_i\\, \\{s_i^2+r_i^2\\, [1-e_i\\, \\cos (2(\\varphi_i-\\phi_i))]\\}^{1\\over 2}+\nr_1^2\\, \\gamma\\, \\cos (2(\\varphi_1-\\varphi_\\gamma))\n\\end{equation}\n\\vspace*{-0.2cm}\n\\normalsize\n\nHere $(r_i,\\varphi_i )$ are polar coordinates with origin at the\ncenter of each galaxy.\n$s$ measures the core radius, $e$ and $\\phi$\nthe ellipticity and position angle of the major axis.\nAt large radius the mass distribution is isothermal, \ngoing as $\\Sigma \\propto r^{-1}$. \nThe lenses will be fixed at $\\vec{x}_{G1}\\!=\\!(0.446,-1.063)''$ and\n$\\vec{x}_{G2}\\!=\\!(-0.276,-0.937)''$, the centroids in H band, which\nare less affected by reddening (Blandford, Surpi \\& Kundi\\'c 1999).\n\nWe minimize a $\\chi^2$ function.\nThe best fit achieved, hereafter Model A, has $\\chi^2\\!=\\!4.0$ and yields \n$H_{\\circ}\\!=\\!100{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$.\nModels with lower values of $H_\\circ$ can also be\nbuilt fixing $H_\\circ$ and fitting the 3 time delays instead of the\ntime delay ratios. As an example we present the results of Model B\nhaving $H_\\circ\\!=\\!69{\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}}$\nand $\\chi^2\\!=\\!24.7$.\nThe parameters and predictions of Model A and B\nare displayed in Tables~\\ref{par} and~\\ref{obs} respectively.\nThey represent reasonable mass distributions given, especially,\nour ignorance of the dark matter distribution (Figure 1).\n\n\\section{Discussion}\n\nA variety of parametrized models \ncan reproduce the point source properties \nof B1608+656. This precludes an accurate determination of H$_\\circ$.\nTo break the degeneracy, extra constraints, associated with\nthe extended emission of the source, have to be incorporated.\nIt is also helpful to specify the distribution of dark matter\non larger scale than the image distribution. A similar conclusion\nhas been drawn by Williams \\& Saha (1999) using pixellated models.\n\\begin{figure}\n\\vspace*{-0.5cm}\n\\begin{center}\n\\leavevmode\n\\epsfysize=1.55in\n\\epsfbox{gsurpi1.eps}\n\\vspace*{-0.2cm}\n\\caption{Surface mass density in Models A and B}\n\\vspace*{-0.5cm}\n\\end{center}\n\\end{figure}\n\n\\begin{table}\n\\caption{Model parameters.} \\label{par}\n\\vspace*{-0.3cm}\n\\begin{center}\n\\footnotesize\n\\begin{tabular}{c|cc|cc}\n\\tableline\n&\\multicolumn{2}{c|}{Model~A~(H$_\\circ\\!=\\!100 \\scriptstyle\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}$)}\n&\\multicolumn{2}{c}{Model~B~(H$_\\circ\\!=\\!69 \\scriptstyle\\rm km\\,s^{\\!-\\!1}Mpc^{\\!-\\!1}$)}\\\\\n%\n\\cline{2-5}\n%\nParameters& G1 & G2 & G1 & G2 \\\\\n%\n\\tableline\n%\ns('') & 0.10 & 0.10 & 0.00 & 0.05 \\\\\nb & 0.9072 & 0.2453 & 0.7797 & 0.3429 \\\\\ne & 0.3269 & 0.6405 & 0.1570 & 0.3149 \\\\\n$\\phi(^\\circ)$ & 163.45 & 154.93 & 172.62 & 160.86 \\\\\n\\tableline\n$\\gamma,\\varphi_{\\gamma}(^\\circ)$ & 0.0876 & -10.92& 0.0473 & 12.47 \\\\\n%\n\\tableline\n%\n\\end{tabular}\n\\end{center}\n\\vspace*{-0.4cm}\n\\caption{Comparison between observations and model predictions.} \\label{obs}\n\\vspace*{-0.3cm}\n\\begin{center}\n\\footnotesize\n\\begin{tabular}{cccc}\n\\tableline\n%\nProperties & \\multicolumn{1}{c}{Observation}\n& \\multicolumn{1}{c}{Model A}\n& \\multicolumn{1}{c}{Model B} \\\\\n%\n\\tableline\n%\n$\\vec x_A('')$ & (~0.0000,~0.0000) $\\pm$ (0.0023,0.0023) & (~0.0000,~0.0000) & (~0.0000,~0.0000)\\\\\n$\\vec x_B('')$ & (-0.7380,-1.9612) $\\pm$ (0.0043,0.0046) & (-0.7382,-1.9613) & (-0.7365,-1.9518)\\\\\n$\\vec x_C('')$ & (-0.7446,-0.4537) $\\pm$ (0.0045,0.0049) & (-0.7443,-0.4544) & (-0.7422,-0.4575)\\\\\n$\\vec x_D('')$ & (~1.1284,-1.2565) $\\pm$ (0.0107,0.0124) & (~1.1271,-1.2582) & (~1.1269,-1.2207)\\\\\n%\n\\tableline\n%\n$F_A / F_B$ & 2.042 $\\pm$ 0.124 & 1.917 & 1.901 \\\\\n$F_C / F_B$ & 1.037 $\\pm$ 0.083 & 1.092 & 1.131 \\\\\n$F_D / F_B$ & 0.350 $\\pm$ 0.055 & 0.428 & 0.504 \\\\\n%\n\\tableline\n%\n%$T_{AB}/T_{CB}$ & 0.79 $\\pm$ 5.0 & 28.4 & 26.4 \\\\ \n%$T_{AB}/T_{DB}$ & 33.0 $\\pm$ 5.0 & 32.0 & 30.8 \\\\\n$T_{AB} (d)$ & 26.0 $\\pm$ 5.0 & 28.4 & 27.6 \\\\ \n$T_{CB} (d)$ & 33.0 $\\pm$ 5.0 & 32.0 & 31.7 \\\\\n$T_{DB} (d)$ & 73.0 $\\pm$ 5.0 & 68.4 & 71.3 \\\\\n%\n\\tableline\n%\n$\\chi^2$ & 0.0 & 4.0 & 24.7 \\\\\n\\tableline\n\\end{tabular}\n\\end{center}\n\\end{table}\n\n\n\\vspace*{-1.0cm}\n\\begin{references}\n\\vspace*{-0.2cm}\n\\small\n\\reference Blandford, R., \\& Kundic, T. 1996 {\\it The Extragalactic\nDistance Scale}, p60\n\\reference Blandford, R., Surpi, G. \\& Kundi\\'c, T. 1999 these proceedings\n\\reference Fassnacht, C. 1999 these proceedings\n\\reference Koopmans, L.V., \\& Fassnacht, C.D. 1999\nastro-ph/9907258, to appear in \\apj\n\\reference Williams, L. L. R. \\& Saha, P. 1999 these proceedings\n\\end{references}\n\\end{document}\n\n\n\n"
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astro-ph0002003
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Compact Stellar Systems in the Fornax Cluster: Super-massive Star Clusters or Extremely Compact Dwarf Galaxies?
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[
{
"author": "M. J. Drinkwater$^{1}$"
},
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"author": "J. B. Jones$^{2}$"
},
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"author": "M. D. Gregg$^{3}$"
},
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"author": "S. Phillipps$^{4}$"
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and %
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[
{
"name": "fss4.tex",
"string": "%\n% LaTeX template file for\n% Publications of the Astronomical Society of Australia.\n% Version 2.4 - 28 April 1997\n%\n%\\documentstyle[12pt,psfig]{article}\n\\documentstyle[11pt,psfig]{article}\n%\n% Baselineskip may be altered if desired.\n%\n\\baselineskip=2em\n%\n% A few definitions. Do not change the reference command.\n%\n\\def\\reference{\\parskip 0pt\\par\\noindent\\hangindent 0.5 truecm}\n\\def\\s{{\\rm\\thinspace s}}\n\\def\\km{{\\rm\\thinspace km}}\n\n\\def\\kms{\\hbox{$\\km\\s^{-1}\\,$}}\n\\def\\bj{\\hbox{$b_j$}} \n\\def\\Bj{\\hbox{$b_j$}} \n\\def\\mo{\\hbox{M$_\\odot$}}\n\\def\\pc{{\\rm\\thinspace pc}}\n\n%\n% Text locations - these may be altered slightly if desired.\n%\n\\textwidth=17.5cm\n\\textheight=24.6 cm\n\\topmargin=-2.5cm\n\\oddsidemargin=-1.0cm\n\\evensidemargin=-1.0cm\n%\n% Start of document\n%\n\\begin{document}\n%\n% Title\n% Capitalise the title normally - do not use ALL CAPS.\n%\n\\title{Compact Stellar Systems in the Fornax Cluster:\nSuper-massive Star Clusters or Extremely Compact Dwarf Galaxies?\n}\n%\n\n% Authors\n% Here comes the author(s) of the paper. Please add the appropriate author\n% names for your paper and indicate within the $^...$ the number(s)\n% which corresponds to the institute(s) of each author. In this example\n% the second author has two institutional affiliations.\n% Add or remove authors as required, maintaining the \\and syntax between\n% each author, but no \\and after the last author.\n% **** IMPORTANT: Leave the closing curly bracket line as is. ******\n\n\\author{M. J. Drinkwater$^{1}$ \\and\n J. B. Jones$^{2}$ \\and\n M. D. Gregg$^{3}$ \\and\n S. Phillipps$^{4}$\n} % IMPORTANT: leave this curly bracket as the first character of this line.\n\n% Date - leave this blank.\n%\\date{resubmitted to PASA, 2000 January 24}\n\\date{to appear in {\\em Publications of the Astronomical Society of Australia}}\n\\maketitle\n\n% Institutions\n% Here fill in your institute name(s) and address(es)\n% The number in $^...$ indicates the author number. For example\n{\\center\n$^1$ School of Physics, University of Melbourne, Victoria 3010,\nAustralia\\\\[email protected]\\\\[3mm]\n$^2$ Department of Physics, University of Bristol, Tyndall Avenue, Bristol, BS8 1TL, England, U.K.\\\\[email protected]\\\\[3mm]\n$^3$ University of California, Davis, \n and Institute for Geophysics and Planetary Physics, \n Lawrence Livermore National Laboratory,\n L-413, Livermore, CA 94550, USA\\\\[email protected]\\\\[3mm]\n$^4$ Department of Physics, University of Bristol, Tyndall Avenue, Bristol, BS8 1TL, England, U.K.\\\\[email protected]\\\\[3mm]\n}\n\n% Abstract\n% Simply place your abstract between the \\begin{abstract} and\n% \\end{abstract} commands.\n%\n\\begin{abstract}\n% Place the abstract here.\n\nWe describe a population of compact objects in the centre of the\nFornax Cluster which were discovered as part of our 2dF Fornax\nSpectroscopic Survey. These objects have spectra typical of old\nstellar systems, but are unresolved on photographic sky survey plates.\nThey have absolute magnitudes $-13<M_B<-11$, so they are 10 times more\nluminous than any Galactic globular clusters, but fainter than any\nknown compact dwarf galaxies. These objects are all within 30\narcminutes of the central galaxy of the cluster, NGC 1399, but are\ndistributed over larger radii than the globular cluster system of that\ngalaxy. \n\nWe suggest that these objects are either super-massive star clusters\n(intra-cluster globular clusters or tidally stripped nuclei of dwarf\ngalaxies) or a new type of low-luminosity compact elliptical dwarf\n(``M32-type'') galaxy. The best way to test these hypotheses will be\nto obtain high resolution imaging and high-dispersion spectroscopy to\ndetermine their structures and mass-to-light ratios. This will allow\nus to compare them to known compact objects and establish if they\nrepresent a new class of hitherto unknown stellar system.\n\\end{abstract}\n\n{\\bf Keywords:}\ngalaxies: star clusters --- galaxies: dwarf --- galaxies: formation\n\n\\bigskip\n\n%\\twocolumn\n\n\\section{Introduction}\n\nIn cold dark matter (CDM) galaxy formation, small dense halos of\ndark matter collapse at high redshift and eventually merge to form the\nlarge virialised galaxy clusters seen today. The CDM model is very\ngood at reproducing large-scale structure, but only very recently have\nthe best numerical simulations (Moore et al.\\ 1998) had the resolution\nto trace the formation of small halos within galaxy clusters, with\nmasses $\\approx 10^9$M$_\\odot$. We do not yet know what the lower mass\nlimit is for the formation of halos in the cluster environment:\ndetermining the lower limit of galaxy mass in clusters will provide an\nimportant constraint on these models. Most of the smallest cluster\ngalaxies are low surface brightness dwarfs for which mass estimates\nare very difficult, though comparison with field low surface\nbrightness dwarfs would suggest that they may be dark matter dominated\n(e.g.\\ Carignan \\& Freeman 1988).\n\nIn this paper we describe a population of small objects we have found\nin the Fornax Cluster (see also Minniti et al.\\ 1998 and Hilker et\nal.\\ 1999) which have high surface brightness. The origin and nature\nof these objects is not yet clear, but if they are a product of the\ngalaxy formation process in clusters, their high surface brightness\nwill make it possible to probe the low-mass limit discussed\nabove. They may represent extreme examples of compact low luminosity\n(``M32-type'') dwarf ellipticals. Alternatively, these objects may be\nsuper-massive star clusters---there is a very large\npopulation of globular clusters associated with the central galaxy of\nthe Fornax Cluster, NGC 1399 (Grillmair et al.\\ 1994). These objects\nare generally similar to Galactic globular clusters with similar\ncolours and luminosities (Forbes et al.\\ 1998). There is evidence that\nthey are not all bound to the NGC 1399 system. Kissler-Patig et al.\\\n(1999) show that the kinematics of 74 of the globular clusters\nindicate that they are associated with the cluster gravitational\npotential rather than that of NGC 1399. They infer that the most\nlikely origin of these globular clusters is that they have been\ntidally stripped from neighbouring galaxies. This has also been\nsuggested by West et al.\\ (1995), although the effect would be diluted\nby the large number of halo stars that would presumably be stripped at\nthe same time.\n\nBassino, Muzzio \\& Rabolli (1994) suggest that the NGC 1399 globular\nclusters are remnants of the nuclei of dwarf nucleated galaxies that\nhave survived the disrupture of being captured by the central cluster\ngalaxy. A related suggestion is a second model proposed by West et\nal.\\ (1995) that intra-cluster globular clusters could have formed in\nsitu in the cluster environment. Bassino et al.\\ (1994) conclude\ntheir discussion by noting that remnant nuclei an order of magnitude\nlarger (and more luminous) than standard globular clusters would also\nbe formed in significant numbers, but that existing globular cluster\nsearches would not have included them. In Section~\\ref{sec-obs} of\nthis paper we describe how the observations of our {\\em Fornax\nSpectroscopic Survey} have sampled this part of the cluster population\nby measuring optical spectra of all objects brighter than $B_J=19.7$\nin the centre of the Fornax Cluster. In Section~\\ref{sec-prop} we\ndescribe the properties of a new population of compact objects found\nin the cluster that appear to be intermediate in size between globular\nclusters and the smallest compact dwarf galaxies. We discuss the\nnature of these objects in Section~\\ref{sec-discuss} and show that\nhigher resolution observations will enable us to determine if they are\nmore like globular clusters or dwarf galaxies.\n\n\\section{Discovery Observations: The Fornax Spectroscopic Survey}\n\\label{sec-obs}\n\nOur {\\em Fornax Spectroscopic Survey}, carried out with the 2dF\nmulti-object spectrograph on the Anglo-Australian Telescope (see\nDrinkwater et al.\\ 2000), is now 87\\% complete in its first field to\na limit of $B_J=19.7$. The 2dF field is a circle of diameter 2 degrees\n(i.e.\\ $\\pi$ square degrees of sky). We have measured optical spectra\nof some 4000 objects (some going fainter than our nominal limit) in a\n2dF field centred on the central galaxy of the Fornax Cluster (NGC\n1399). This survey is unique in that the targets (selected from\ndigitised UK Schmidt Telescope photographic sky survey plates) include\n{\\em all} objects, both unresolved (``stars'') and resolved\n(``galaxies'') in this large area of sky. The resolved objects\nmeasured are mostly background galaxies as expected with a minor\ncontribution from Fornax Cluster members. The unresolved objects are\nmostly Galactic stars and distant AGN, also as expected, but some are\ncompact starburst (and post-starburst) galaxies beyond the Fornax\nCluster (Drinkwater et al 1999a).\n\nFinally, in addition to the dwarf galaxies already listed in the\nFornax Cluster Catalog (FCC: Ferguson 1989) which we have confirmed as\ncluster members, we have found a sample of five very compact objects\nat the cluster redshift which are unresolved on photographic sky\nsurvey plates and not included in the FCC. These new members of the\nFornax Cluster are listed in Table~\\ref{tab-list} along with their\nphotometry measured from the UKST plates.\n\n\n\\begin{table*}\n\\caption{The new compact objects\n\\label{tab-list}}\n\\center\n\\begin{tabular}{lllll}\n\\hline\nName & RA (J2000) Dec & $B_J$& $M_B$ & cz \\\\\n & & (mag)& (mag) & (\\kms) \\\\\n\\hline\nThales 1 & 03 37 3.30 -35 38 4.6 &19.85 & $-11.1$ & 1507 \\\\\nThales 2 & 03 38 6.33 -35 28 58.8 &18.85 & $-12.1$ & 1328 \\\\\nThales 3$^1$ & 03 38 54.10 -35 33 33.6 &17.68 & $-13.2$ & 1595 \\\\\nThales 4$^2$ & 03 39 35.95 -35 28 24.5 &18.82 & $-12.1$ & 1936 \\\\\nThales 5 & 03 39 52.58 -35 04 24.1 &19.66 & $-11.2$ & 1337 \\\\\n\\hline\n\\end{tabular}\n\nNotes: (1) CGF 1-4 (2) CGF 5-4, both in Hilker et al.\\ (1999) \n\n\\end{table*}\n\n\nOur 2dF measurements of unresolved objects are 80\\% complete in the\nmagnitude range of these objects ($17.5<\\bj<20.0$). There is therefore\nabout one more similar compact object still to be found in in our\ncentral 2dF field. The number density of these objects is\ntherefore $6\\pm3$ per 2dF field ($\\pi$ square degrees). Two of\nthe objects (the two brightest) were also identified as cluster\nmembers by Hilker et al.\\ (1999). Hilker et al.\\ measured spectra of\nabout 50 galaxies brighter than $V=20$ in a square region of width\n0.25 degrees at the centre of the Fornax Cluster. In the ``galaxies''\nthey include objects which were very compact, but still resolved. By\ncontrast our own survey covers a much larger area and also includes all\nunresolved objects.\n\n\\section{Properties of the compact objects}\n\\label{sec-prop}\n\n\\subsection{Sizes}\n\nThese object images are unresolved and classified ``stellar'' in our\nUKST plate data, although imaging with the CTIO Curtis Schmidt shows\nthat the brightest two objects have marginal signs of extended\nstructure. In Fig.~\\ref{fig-image} we present R-band (Tech Pan\nemulsion + OG 590 filter) photographic images of these compact objects\nfrom the UKST. These were taken in seeing of about 2.2 arcseconds FWHM\nand the third object (Thales\\footnote{Thales of Miletus was the first\nknown Greek philosopher and scientist and possibly the earliest\nastronomer.} 3) is resolved with a 3.2 arcsecond FWHM. Applying a\nvery simple deconvolution of the seeing this corresponds to a scale\nsize (HWHM) of about 80 pc. This is much larger than any known\nglobular cluster, so this object, at least, is not a globular cluster.\nThe other objects are all unresolved, so must have scale sizes smaller\nthan this.\n\n\\begin{figure*}\n\\hfil \\psfig{file=fig_im.eps,angle=0,width=18cm}\n\\caption{R-band photographic images of the new compact objects. The\nimages are all from a single UKST exposure on Tech-Pan emulsion,\ndigitised by SuperCOSMOS (Miller et al.\\ 1992). Each image is 2.5\narcminutes across with North at the top and East to the left.\n\\label{fig-image}}\n\\end{figure*}\n\n\\subsection{Luminosity and Colours}\n\nThese new objects have absolute magnitudes $-13<M_B<-11$, based on a\ndistance modulus of 30.9 mag to the Fornax Cluster (Bureau et al.\\\n1996). These values are at the lower limit of dwarf galaxy\nluminosities (Mateo 1998), but are much more luminous than any known\nGalactic globular clusters (Harris, 1996) and the most luminous of the\nNGC1399 globulars (Forbes et al.\\ 1998) which have $M_B\\approx-10$.\nThe luminosities of the compact objects are compared to several other\npopulations of dwarf galaxy and star cluster in Fig.~\\ref{fig-lf}. We\nnote that the magnitude limit of the 2dF data corresponds to an\nabsolute magnitude of $M_B\\approx-11$ here. In order of decreasing\nluminosity the first comparison is with the dwarf ellipticals listed\nin the FCC as members of the Fornax Cluster. The possible M32-type\ngalaxies in the FCC are not included as none of them have yet been\nshown to be cluster members (Drinkwater, Gregg \\& Holman 1997). The\nFigure shows that the Fornax dEs have considerable overlap in\nluminosity with the compact objects, but morphologically they are very\ndifferent, being fully resolved low surface brightness\ngalaxies. Recently, several new compact dwarf galaxies have been\ndiscovered in the Fornax Cluster (Drinkwater \\& Gregg 1998) but these\nare all brighter than $M_B=-14$ and do not match any of the objects we\ndiscuss here. Binggeli \\& Cameron (1991) measured the luminosity\nfunction of the nuclei of nucleated dwarf elliptical galaxies in the\nVirgo Cluster. The Figure shows that this also overlaps the\ndistribution of the new compact objects. In this case the morphology\nis the same, so the compact objects could originate from the dwarf\nnuclei. The Figure also shows the luminosity functions of both the NGC\n1399 globular clusters (Bridges, Hanes \\& Harris, 1991) and Galactic\nglobular clusters (Harris 1996). These are quite similar and have no\noverlap with the compact objects.\n\nFor completeness we note that the luminosities of the compact objects\nhave considerable overlap with the luminosities of dwarf galaxies in\nthe Local Group (Mateo 1998), but even the most compact of the Local\nGroup dwarfs, Leo I ($M_B=-11.1$) would be resolved ($r_e\\approx 3''$)\nin our images at the distance of Fornax. The only population they\nmatch in both luminosity and morphology is the bright end of the\nnuclei of nucleated dwarf ellipticals.\n\n\\begin{figure}\n\\hfil \\psfig{file=fig_lf.eps,angle=0,width=9cm}\n\\caption{Distribution of absolute magnitude of the compact objects\n(filled histogram) compared to dEs in the Fornax Cluster (Ferguson\n1989; solid histogram), the nuclei of dE,Ns in the Virgo Cluster\n(Binggeli \\& Cameron 1991; short dashes), a model fit to the globular\nclusters around NGC 1399 (Bridges, Hanes \\& Harris, 1991; long dashes)\nand Galactic globular clusters (Harris 1996; dotted). Note: the\nmagnitude limit of our survey that found the compact objects\ncorresponds to $M_B=-11$.\n\\label{fig-lf}}\n\\end{figure}\n\n\n\\subsection{Spectral Properties}\n\nThe 2dF discovery spectra of these compact objects are shown in\nFigure~\\ref{fig-spec}. They have spectra similar to those of early\ntype dwarf galaxies in the sample (two are shown for comparison in the\nFigure) with no detectable emission lines. As part of the spectral\nidentification process in the {\\em Fornax Spectroscopic Survey}, we\ncross-correlate all spectra with a sample of stellar templates from\nthe Jacoby et al.\\ (1984) library. The spectra of the new compact\nobjects were best fit by K-type stellar templates, consistent with an\nold (metal-rich) stellar population. The dE galaxies observed with the\nsame system by contrast are best fit by younger F and early G-type\ntemplates. This gives some indication in favour of the compact objects\nbeing related to globular clusters, although we note that two of them\nwere analysed by Hilker et al.\\ (1999) in more detail without any\nconclusive results. We do not have the spectrum of a dE nucleus\navailable for direct comparison, but since our 2dF spectra are taken\nthrough a 2 arcsec diameter fibre aperture, the spectrum of FCC 211, a\nnucleated dE, is dominated by the nucleus. This spectrum was fitted by\na younger F-type stellar template, again suggestive of a younger\npopulation than the compact objects. We cannot draw any strong\nconclusions from these low-resolution, low signal-to-noise spectra.\n\n\\begin{figure}\n%\\hfil \\psfig{file=plot_spec.eps,angle=0,width=11.2cm}\n\\hfil \\psfig{file=plot_spec.eps,angle=0,width=9cm}\n\\caption{2dF discovery spectra of the five compact objects as well as\ntwo cluster dwarf galaxies for comparison. Note: the large scale\nripple in the spectrum of Thales~5 is an instrumental effect caused by\ndeterioration in the optical fibre used.\n\\label{fig-spec}}\n\\end{figure}\n\n\\subsection{Radial Distribution}\n\nThe main advantage of our survey over the previous studies of the NGC\n1399 globular cluster system (e.g.\\ Grillmair et al.\\ 1994, Hilker et\nal.\\ 1999) is that we have complete spectroscopic data over a much\nlarger field, extending to a radius of 1 degree (projected distance of\n270 kpc) from the cluster centre. This means that we can determine the\nspatial distribution of the new compact objects. In\nFigure~\\ref{fig-radial} we plot the normalised, cumulative radial\ndistribution of the new compact objects compared to that of foreground\nstars and cluster galaxies. This plot is the one used to calculate\nKolmogorov-Smirnov statistics and allows us to compare the\ndistributions of objects independent of their mean surface\ndensities. It is clear from the Figure that the new compact objects\nare very concentrated towards the centre of the cluster, at radii\nbetween 5 and 30 arcminutes (20--130 kpc). Their distribution is more\ncentrally concentrated than the King profile fitted to cluster members\nby Ferguson (1989) with a core radius of 0.7 degrees (190 kpc). The\nKolmogorov-Smirnov (KS) test gives a probability of 0.01 that the\ncompact objects have the same distributions as the FCC galaxies: they\nare clearly not formed (or acreted) the same way as average cluster\ngalaxies. To test the hypothesis that the compact objects are formed\nfrom nucleated dwarfs, we also plot the distribution of all the FCC\nnucleated dwarfs, as these are more clustered than other dwarfs\n(Ferguson \\& Binggeli 1994). However in the central region of interest\nhere the nucleated dwarf profile lies very close to the King profile\nof all the FCC galaxies, so this does not provide any evidence for a\ndirect link with the new compact objects.\n\nWest et al.\\ (1995) suggest that a smaller core radius should be used\nfor intra-cluster globular clusters (GCs). This profile, also shown in\nthe Figure, is more consistent with the distribution of the new\nobjects: the KS probability of the compact objects being drawn from\nthis distribution is 0.39.\n\n% calculate the KS statistics:\n% West: 0.8993707895 0.3936000764\n% FCC: 1.594319463 0.01239375863\n\n\n\\begin{figure}\n\\hfil \\psfig{file=fig_radial.eps,angle=0,width=9cm}\n\\caption{Cumulative radial distribution of the new compact objects\ncompared to the predicted distribution for intra-cluster globular\nclusters (West et al 1995) and the profile fit to the distribution of all\nFCC. Also shown is the distribution of all nucleated dwarfs in the FCC\nand all the unresolved objects (stars) observed in our 2dF survey.\n\\label{fig-radial}}\n\\end{figure}\n\nWe also note that the radial distribution of the compact objects is\nmuch more extended than the NGC 1399 globular cluster system as\ndiscussed by Grillmair et al.\\ and extends to three times the\nprojected radius of that sample. It unlikely that all the compact\nobjects are associated with NGC 1399. This is emphasised by a finding\nchart for the central 55 arcminutes of the cluster in\nFig.~\\ref{fig-apm} which indicates the location of the compact\nobjects. They are widely distributed over this field and Thales~3 in\nparticular is much closer to NGC 1404, although we note that its\nvelocity is not close to that of NGC 1404 (see below).\n\n\\begin{figure*}\n\\hfil \\psfig{file=figure5_10n.eps,width=15cm}\n\\caption{The central region of the Fornax Cluster with the positions\nof the new compact objects indicated by squares. This R-band\nphotographic image is from a single UKST exposure on Tech-Pan\nemulsion, digitised by SuperCOSMOS (Miller et al.\\ 1992).\n\\label{fig-apm}}\n\\end{figure*}\n\n\\subsection{Velocity Distribution}\n\nWe have some limited information from the radial velocities of the\ncompact objects. The mean velocity of all 5 ($1530\\pm110 \\kms$) is\nconsistent with that of the whole cluster ($1540\\pm50 \\kms$) (Jones \\&\nJones (1980). However, given the small sample, it is also consistent\nwith the velocity of NGC 1399 ($1425\\pm 4 \\kms$) as might be expected\nfor a system of globular clusters. Interestingly, the analysis of the\ndynamics of 74 globular clusters associated with NGC 1399 by\nKissler-Patig et al.\\ (1999) notes that their radial velocity\ndistribution has two peaks, at about 1300 and 1800\\kms.\nOur sample is far too small to make any conclusions about the dynamics\nof these objects at present.\n\n\\section{Discussion}\n\\label{sec-discuss}\n\nWe cannot say much more about the nature of these objects on the basis\nof our existing data. In ground-based imaging, they are intermediate\nbetween large GCs and small compact dwarf galaxies, so it becomes\nalmost a matter of semantics to describe them as one or the other. The\nmost promising way to distinguish between these possibilities is to\nmeasure their mass-to-light (M/L) ratios. If they are large, but\notherwise normal, GCs, they will be composed entirely of stars giving\nvery low M/L. If they are the stripped nuclei of dwarf galaxies we\nmight expect them to be associated with some kind of dark halo, but we\nwould not detect the dark halos at the small radii of these nuclei, so\nwe would also measure small M/L values. Alternatively, these objects\nmay represent a new, extreme class of compact dwarf elliptical\n(``M32-type'') galaxy. These would presumably have formed by\ngravitational collapse within dark-matter halos, so would have high\nmass-to-light-ratios, like dwarf galaxies in the Local Group (Mateo\n1998). One argument against this interpretation is the apparent lack\nof M32-like galaxies at brighter luminosities (Drinkwater \\& Gregg\n1998). If the compact objects are dwarf galaxies, they will represent\nthe faintest M32-like galaxies ever found. They may also fill in the\ngap between globular clusters and the fainter compact galaxies in the\nsurface brightness vs.\\ magnitude distribution given by Ferguson \\&\nBinggeli (1994).\n\nA further possibility is that these are small scale length ($\\sim\n100$~pc) dwarf spheroidal galaxies of only moderately low surface\nbrightness. While Local Group dSphs of equivalent luminosities\ngenerally have substantially larger scale sizes (and consequently\nlower surface brightnesses) (Mateo, 1998), Leo I for example has $M_B\n= -11.0$, and a scale length of only 110~pc (Caldwell et al., 1992),\nbut as we discuss above this would be resolved in our existing\nimaging.\n\nOur existing data will only allow us to estimate a conservative upper\nlimit to the mass of these objects. If we say that the core radii of\nthe objects are less than 75\\pc\\ and the velocity dispersions are less\nthan 400\\kms\\ (the resolution of our 2dF spectra) we find that the\nvirial mass must be less than $10^{10}\\mo$. For a typical luminosity\nof $M_B=-12$ this implies that $M/L < 2\\times 10^{3}$. This is not a\nvery interesting limit, so we plan to reobserve these objects at\nhigher spectral resolution from the ground and higher spatial\nresolution with the {\\em Hubble Space Telescope} (HST) in order to be\nsensitive to $M/L\\approx 100$. This will allow us to distinguish\nglobular clusters from dwarf galaxies.\n\nIn order to demonstrate what we could measure with high-resolution\nimages, we present two extreme possibilities in\nFigure~\\ref{fig-profile}: a very compact Galactic globular cluster and\na dwarf galaxy with an $r^{1/4}$ profile ($r_e=0.2$ arcsec), both\nnormalised to magnitudes of $B=19$ ($V=18.4$) and the Fornax cluster\ndistance. We also plot the PSF of the {\\em Space Telescope Imaging\nSpectrograph} (STIS) in the Figure for reference. The globular cluster\nprofile is that of NGC 2808 (Illingworth \\& Illingworth 1976) with the\nradius scaled to the distance of the Fornax Cluster and the surface\nbrightness then scaled to give the desired apparent magnitude. The\nglobular cluster profile is very compact and will only just be\nresolved with HST, but it will clearly be differentiated from the\ndwarf galaxy profile.\n\n\\begin{figure}\n\\hfil \\psfig{file=hst_glob_f2.eps,angle=0,width=9cm}\n\\caption{Predicted radial surface brightness profiles of the compact\nobjects in two extreme cases: (A) a Galactic globular cluster\n(Illingworth \\& Illingworth 1976) scaled 3 mag brighter in surface\nbrightness, and (B) a compact dwarf galaxy with an $r^{1/4}$\nprofile. Both are scaled to have total magnitudes of B=19 mag; they\nare not corrected for instrumental PSF which is also shown\n(C). \\label{fig-profile}}\n\\end{figure}\n\nIn addition to measuring the size of these objects for the mass\nmeasurement, the radial surface brightness profiles may also give\ndirect evidence for their origin and relationship to other kinds of\nstellar systems. For example, if they are the stripped nuclei of\ngalaxies, the remnants of the outer envelope might show up in the HST\nimages as an inflection in the surface brightness profile at large\nradius.\n\n\\section{Summary}\n\nWe have reviewed the observed properties of these new compact objects\ndiscovered in the Fornax Cluster. Their luminosities are intermediate\nbetween those of known globular clusters and compact dwarf galaxies,\nbut they are consistent with the bright end of the luminosity function\nof the the nuclei of nucleated dwarf ellipticals. The 2dF spectra are\nsuggestive of old (metal-rich) stellar populations, more like globular\nclusters than dwarf galaxies. Finally the radial distribution of the\ncompact objects is more centrally concentrated than cluster galaxies\nin general, but extends further than the known globular cluster system\nof NGC 1399.\n\nThese objects are most likely either massive star clusters (extreme\nglobular clusters or tidally-stripped dwarf galaxy nuclei) or very\ncompact, low-luminosity dwarf galaxies. In the latter case these new\ncompact objects would be very low-luminosity counterparts to the\npeculiar compact galaxy M32. This would be particularly interesting\ngiven the lack of M32-like galaxies at brighter luminosities\n(Drinkwater \\& Gregg 1998). With higher resolution images and spectra\nwe will be able to measure the mass-to-light ratios of these objects\nand determine which of these alternatives is correct.\n\n\n\n\\section*{Acknowledgements}\n\nWe thank the referee for helpful suggestions which have improved the\npresentation of this work. We wish to thank Dr.\\ Harry Ferguson for\nhelpful discussions and for providing the STIS profile. We also\nthank Dr.\\ Trevor Hales for assistance in the naming of the\nobjects. MJD acknowledges support from an Australian Research Council\nLarge Grant.\n\n\\section*{References}\n\n\\reference Caldwell, N., Armandroff, T.E., Seitzer, P., Da Costa,\nG.S., 1992, AJ, 103, 840\n\\reference Bassino, L.P., Muzzio, J.C., Rabolli, M. 1994, ApJ, 431, 634\n\\reference Binggeli, B., Cameron, L.M., 1991, A\\&A, 252, 27\n\\reference Bridges, T.J., Hanes, D.A., Harris, W.E., 1991, AJ, 101, 469\n\\reference Bureau, M., Mould, J.R., Staveley-Smith, L., 1996, ApJ,\n463, 60\n\\reference Carignan, C., Freeman, K.C. 1988, ApJ, 332, L33\n\\reference Drinkwater, M.J., Gregg, M.D., 1998, MNRAS, 296, L15\n\\reference Drinkwater, M.J., Gregg, M.D., Holman, B.A., 1997 in\nArnaboldi M., Da Costa G.S., Saha P., eds, ASP Conf. Ser. Vol. 116,\nThe Second Stromlo Symposium: The Nature of Elliptical Galaxies.\nAstron. Soc. Pac., San Francisco, p. 287\n\\reference Drinkwater, M.J., Phillipps, S., Gregg, M.D., Parker, Q.A.,\nSmith, R.M., Davies, J.I., Jones, J.B., Sadler, E.M., 1999a, ApJ, 511, L97\n\\reference Drinkwater, M.J., Phillipps, S., Jones, J.B., Gregg, M.D.,\nDeady, J.H., Davies, J.I., Parker, Q.A., Sadler, E.M., \nSmith, R.M. 2000, A\\&A, submitted\n\\reference Ferguson H.C., 1989, AJ, 98, 367\n\\reference Ferguson H.C., Binggeli, B., 1994, A\\&ARv, 6, 67\n\\reference Forbes, D.A., Grillmair, C.J., Williger, G.M., Elson,\nR.A.W., Brodie, J.P. 1998, MNRAS, 293, 325\n\\reference Grillmair, C.J., Freeman, K.C., Bicknell, G.V., Carter, D.,\nCouch, W.J., Sommer-Larsen, J., Taylor, K. 1994, ApJ, 422, L9\n\\reference Harris, W.E., 1996, AJ, 112, 1487\n\\reference Hilker, M., Infante, L., Vieira, G., Kissler-Patig, M.,\nRichtler, T., 1999, A\\&AS, 134, 75\n\\reference Illingworth, G., Illingworth, W. 1976, ApJSup, 30, 227\n\\reference Jacoby, G.H., Hunter, D.A., Christian, C.A., 1984, ApJSup,\n56, 257\n\\reference Jones, J.E., Jones, B.J.T. 1980, MNRAS, 191, 685\n\\reference Kissler-Patig, M., Grillmair, C.J., Meylan, G., Brodie,\nJ.P., Minniti, D., Goudfrooij, P., 1999, AJ, 117, 1206\n\\reference Mateo, M., 1998, Ann. Rev. Astron. Astrophys., 36, 435\n\\reference Minniti, D., Kissler-Patig, M., Goudfrooij, P., Meylan, G., 1998, AJ, 115, 121\n\\reference Miller, L. A., Cormack, W., Paterson, M., Beard, S., Lawrence, L., \n1992, in `Digitised Optical Sky Surveys', eds. H.T. MacGillivray, \nE.B Thomson, Kluwer Academic Publishers, p. 133\n\\reference Moore, B., Governato, F., Quinn, T., Stadel, J., Lake, G. 1998, ApJ, 499, L5\n\\reference West, M.J., Cote, P., Jones, C., Forman, W., Marzke, R.O. 1995 ApJ 453 L77\n\n\\end{document}\n\n"
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astro-ph0002004
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Extracting Energy from a Black Hole through Its Disk
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"author": "Li-Xin Li"
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When some magnetic field lines connect a Kerr black hole with a disk rotating around it, energy and angular momentum are transferred between them. If the black hole rotates faster than the disk, $ca/GM_H>0.36$ for a thin Keplerian disk, then energy and angular momentum are extracted from the black hole and transferred to the disk ($M_H$ is the mass and $a M_H$ is the angular momentum of the black hole). This way the energy originating in the black hole may be radiated away by the disk. The total amount of energy that can be extracted from the black hole spun down from $ca/GM_H = 0.998$ to $ca/GM_H = 0.36$ by a thin Keplerian disk is $\approx 0.15 M_Hc^2$. This is larger than $\approx 0.09 M_Hc^2$ which can be extracted by the Blandford-Znajek mechanism.
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"name": "bh_disk.tex",
"string": "% bh_disk.tex for ApJ Letters\n\n\\documentstyle[12pt,aaspp4]{article}\n\\begin{document}\n\n\\title{Extracting Energy from a Black Hole through Its Disk}\n\n\\author{Li-Xin Li}\n\\affil{Princeton University Observatory, Princeton, NJ 08544--1001, USA}\n\\affil{E-mail: [email protected]}\n%\\affil{(November 5, 1999; Revised January 14, 2000)}\n\n\\begin{abstract}\nWhen some magnetic field lines connect a Kerr black hole with a disk \nrotating around it, energy and angular momentum are transferred between them. \nIf the black hole rotates faster than the disk, $ca/GM_H>0.36$ for a thin \nKeplerian disk, then energy and angular momentum are extracted from the black \nhole and transferred to the disk ($M_H$ is the mass and $a M_H$ is the \nangular momentum of the black hole). This way the energy originating\nin the black hole may be radiated away by the disk.\n\nThe total amount of energy that can be extracted from the black hole\nspun down from $ca/GM_H = 0.998$ to $ca/GM_H = 0.36$ by a thin Keplerian\ndisk is $\\approx 0.15 M_Hc^2$. This is larger than $\\approx 0.09 M_Hc^2$ \nwhich can be extracted by the Blandford-Znajek mechanism.\n\\end{abstract}\n\n\\keywords{black hole physics --- accretion disks --- magnetic fields}\n\n%\\section 1\n\\section{Introduction}\nExtraction of energy from a black hole or an accretion disk through\nmagnetic braking has been investigated by many people. As a rotating black hole\nis threaded by magnetic field lines which connect with remote astrophysical loads, \nenergy and angular momentum\nare extracted from the black hole and transported to the remote loads via Poynting\nflux (Blandford \\& Znajek 1977; Macdonald \\& Thorne 1982; Phinney 1983). This is\nusually called the Blandford-Znajek mechanism and has been suggested to be\na plausible process for powering jets in active galactic nuclei (Rees, Begelman,\nBlandford, \\& Phinney 1982; Begelman, Blandford, \\& Rees 1984) and gamma ray\nbursts (Paczy\\'nski 1993; Lee, Wijers, \\& Brown 1999). Similar\nprocess can happen to an accretion disk when some of magnetic field\nlines threading the disk are open and connect with remote astrophysical loads\n(Blandford 1976; Blandford \\& Znajek 1977; Macdonald \\& Thorne 1982; Livio,\nOgilvie, \\& Pringle 1999; Li 1999).\n\nIn this paper we investigate the effects of magnetic field lines connecting a\nKerr black hole with a disk surrounding it. This kind of\nmagnetic field lines are expected to exist and have important effects (Macdonald\n\\& Thorne 1982; Blandford 1999, 2000; Gruzinov 1999). We find that, with the\nexistence of such magnetic coupling between the black hole and the disk, energy\nand angular momentum are transfered between them. If the black hole rotates faster \nthan the disk, energy and angular momentum are extracted from the black hole and \ntransferred to the disk via Poynting flux. This is the case when\n$a/M_H>0.36$ for a thin Keplerian disk, where $M_H$ is the mass of the \nblack hole and $aM_H$ is the angular momentum of the black hole. \nThroughout the paper we use the geometric units with $G = c = 1$. The energy \ndeposited into the disk by the black hole is eventually radiated to infinity\nby the disk. This provides a way for extracting energy from a black hole through its\ndisk. If the disk has no accretion (or the accretion rate is very low), \nthe power of the disk \nessentially comes from the rotational energy of the black hole. We will show\nthat the magnetic coupling between the black hole and the disk has a higher\nefficiency in extracting energy from a Kerr black hole than the Blandford-Znajek\nmechanism.\n\n\n%\\section 2\n\\section{Transfer of Energy and Angular Momentum between a Black Hole and Its\nDisk by Magnetic Coupling}\nSuppose a bunch of magnetic field lines connect a rotating black\nhole with a disk surrounding it. Due to the rotation of the black hole\nand the disk, electromotive forces are induced on both the black hole's horizon \nand the disk (Macdonald \\& Thorne 1982; Li 1999)\n\\begin{eqnarray}\n {\\cal E}_H = {1\\over 2\\pi}\\Omega_H \\Delta\\Psi\\,, \\hspace{1cm}\n {\\cal E}_D = -{1\\over 2\\pi}\\Omega_D \\Delta\\Psi\\,,\n \\label{emf}\n\\end{eqnarray}\nwhere $\\Omega_H$ is the angular velocity of the black hole, $\\Omega_D$ is the\nangular velocity of the disk, $\\Delta\\Psi$ is the magnetic flux connecting\nthe black hole with the disk. The black hole and the disk form a closed\nelectric circuit, the electric current flows through the magnetic field lines \nconnecting them. Suppose the disk and the black hole rotates in the same direction, \nthen ${\\cal E}_H$ and ${\\cal E}_D$ have opposite signs. This means that energy\nand angular momentum are transferred\neither from the black hole to the disk or from the disk to the black hole, the\ndirection of transfer is determined by the sign of ${\\cal E}_H + {\\cal E}_D$. By the\nOhm's law, the current is $I = ({\\cal E}_H+{\\cal E}_D)/Z_H =\n\\Delta\\Psi(\\Omega_H-\\Omega_D)/(2\\pi Z_H)$, where $Z_H$ is the resistance of the \nblack hole which is of several hundred Ohms (the disk is perfectly conducting so\nits resistance is zero). The power deposited into the disk by the black hole is\n\\begin{eqnarray}\n P_{HD} = - I {\\cal E}_D\n = \\left({\\Delta\\Psi\\over 2\\pi}\\right)^2 \\,{\\Omega_D\n \\left(\\Omega_H - \\Omega_D\\right)\n \\over Z_H}\\,.\n \\label{pow3}\n\\end{eqnarray}\nThe torque on the disk produced by the black hole is\n\\begin{eqnarray}\n T_{HD} = {I\\over 2\\pi}\\Delta\\Psi =\n \\left({\\Delta\\Psi\\over 2\\pi}\\right)^2 \\,{\\left(\\Omega_H -\n\t \\Omega_D\\right)\n \\over Z_H}\\,.\n \\label{toq}\n\\end{eqnarray}\nAs expected, we have $P_{BH} = T_{BH}\\Omega_D$.\n\nThe signs of $P_{HD}$ and $T_{HD}$ are determined by the sign of \n$\\Omega_H-\\Omega_D$. When $\\Omega_H > \\Omega_D$, we have $P_{HD}>0$ \nand $T_{HD}>0$, energy and angular\nmomentum are transferred from the black hole to the disk.\nWhen $\\Omega_H < \\Omega_D$, we have $P_{HD} < 0$ and $T_{HD}<0$, energy \nand angular momentum are transferred from the disk to the black hole so the\nblack hole is spun up. For a disk with non-rigid rotation, $\\Omega_D$\nvaries with radius. For fixed values of $\\Delta\\Psi$,\n$\\Omega_H$, and $Z_H$, $P_{HD}$ peaks at $\\Omega_D = \\Omega_H/2$.\nHowever for realistic cases which is most important is when the\nmagnetic field lines touch the disk close to the inner boundary,\nso $\\Omega_D$ in Eq.~(\\ref{pow3}) and Eq.~(\\ref{toq}) can be taken \nto be the value\nat the inner boundary of the disk. According to Gruzinov (1999) the\nmagnetic fields will be more unstable against screw instability if the \nfoot-points \nof the field lines on the disk are far from the inner boundary of the disk.\n\nFor a thin Keplerian disk around a Kerr black hole in the equatorial plane, the\nangular velocity of the disk is (Novikov \\& Thorne 1973)\n\\begin{eqnarray}\n \\Omega_D(r)=\\left({M_H\\over r^3}\\right)^{1/2}{1\\over 1+a\\left(M_H/r^3\n \\right)^{1/2}}\\,,\n \\label{wd}\n\\end{eqnarray}\nwhere $r$ is the Boyer-Lindquist radius in Kerr spacetime. $\\Omega_D(r)$\ndecreases with increasing $r$. The angular velocity of a Kerr black hole is\n\\begin{eqnarray}\n \\Omega_H = {a\\over 2M_H r_H}\\,,\n \\label{wh}\n\\end{eqnarray}\nwhere $r_H = M_H + \\sqrt{M_H^2-a^2}$ is the radius of the event horizon.\n$\\Omega_H$ is constant on the horizon. The inner boundary of a Keplerian disk\nis usually assumed to be at the marginally stable orbit with radius\n(Novikov \\& Thorne 1973)\n\\begin{eqnarray}\n r_{ms}=M_H\\left\\{3+z_2-\\left[(3-z_1)(3+z_1+2z_2)\\right]^{1/2}\\right\\}\\,,\n \\label{rms}\n\\end{eqnarray}\nwhere\n\\begin{eqnarray}\n z_1=1+\\left(1-a^2/M_H^2\\right)^{1/3}\\left[\\left(1+a/M_H\\right)^{1/3}+\n \\left(1-a/M_H\\right)^{1/3}\\right]\\,,\n \\label{rms2}\n\\end{eqnarray}\nand\n\\begin{eqnarray}\n z_2=\\left(3a^2/M_H^2+z_1^2\\right)^{1/2}\\,.\n \\label{rms3}\n\\end{eqnarray}\nInserting Eq.~(\\ref{rms}) into\nEq.~(\\ref{wd}), we obtain the angular velocity of the disk at its inner \nboundary:\n$\\Omega_{ms} = \\Omega_D(r_{ms})$. For the Schwarzschild case (i.e $a = 0$) we \nhave $r_{ms} = 6 M_H$ and $\\Omega_{ms} = 6^{-3/2}M_H^{-1}\\equiv\\Omega_0$.\n\nAssuming the magnetic field lines touch the disk close to the inner boundary, \nwe have $P_{HD} \\approx P_0 f$\nwhere\n\\begin{eqnarray}\n P_0 = \\left({\\Delta\\Psi\\over 2\\pi}\\right)^2 {\\Omega_0^2\\over Z_H}\n \\label{psch}\n\\end{eqnarray}\nis the value of $-P_{HD}$ for the Schwarzschild case, and\n\\begin{eqnarray}\n f = {\\Omega_{ms}\\left(\\Omega_H -\\Omega_{ms}\\right)\\over \\Omega_0^2}\n \\label{ratio}\n\\end{eqnarray}\nis a function of $a/M_H$ only. The variation of $P_{HD}$ with $a/M_H$\nis shown in Fig.~\\ref{figure1}. We see that $P_{HD}>0$ when $0.36 < a/M_H <1$,\n$P_{HD}<0$ when $0\\le a/M_H<0.36$. $P_{HD}=0$ at $a/M_H \\approx 0.36$ and\n$a/M_H=1$ since $P_{HD}\\propto \\Omega_H-\\Omega_{ms}$ and $\\Omega_H = \\Omega_{ms}$ \nwhen $a/M_H\\approx 0.36$ and $a/M_H =1$.\nFor fixed $\\Delta\\Psi$, $M_H$, and $Z_H$, $P_{HD}$ peaks\nat $a/M_H\\approx 0.981$. $T_{HD}$ always has the same sign as $P_{HD}$ since\n$P_{HD} = T_{HD}\\Omega_D$ for a perfectly conducting disk.\n\n\n%\\section 3\n\\section{Extracting Energy from a Black Hole through Its Disk}\nWhen $a/M_H>0.36$, energy and angular momentum are extracted from the black hole\nand transferred to the disk. So a fast rotating black hole can pump its rotational\nenergy into a disk surrounding it through magnetic coupling between them.\nOnce the energy gets into the disk, it can be radiated to infinity either\nin the form of Poynting flux associated with jets or winds, or in the form\nof thermal radiation associated with dissipative processes in the disk.\nIf the disk is not accreting or its accretion rate is very low,\nthen the disk's power \ncomes from the rotational energy of the black hole.\nThis provides a way for {\\em indirectly} extracting energy from a rotating \nblack hole. Note, that the Blandford-Znajek mechanism is a way for \n{\\em directly} extracting energy from a rotating black hole\nto the remote load.\n\nIt is possible that the Blandford-Znajek mechanism provides a very ``clean''\nenergy beam, while energy extracted from the disk is ``dirty'', contaminated\nby matter from the disk corona (R. D. Blandford 1999a, private communication).\nHowever, we must keep in mind that there exists no quantitative model \ndemonstrating how to generate clean energy with the Blandford-Znajek process.\n\nLet us consider again our case, in which\nKerr black hole loses its energy and angular momentum through the \nmagnetic interaction with a thin Keplerian disk, with the magnetic field lines \ntouching the disk close to the marginally stable orbit. \nThe evolution of the mass \nand angular momentum of the black hole are given by\n\\begin{eqnarray}\n {d M_H\\over dt} = -2 P_{HD}\\,,\n \\hspace{1cm} {d J_H\\over dt} = -2 T_{HD}\\,,\n \\label{evol}\n\\end{eqnarray}\nwhere $P_{HD}$ and $T_{HD}$ are given by Eq.~(\\ref{pow3}) and Eq.~(\\ref{toq})\nrespectively, the factors $2$ come from the fact that a disk has\ntwo faces. From Eq.~(\\ref{evol}) we obtain ${dJ_H/ dM_H} = {1/\\Omega_{ms}}$,\nwhere we have used $P_{HD} \\approx T_{HD}\\Omega_{ms}$. Define the spin of a Kerr\nblack hole by $s \\equiv a/M_H = J_H/M_H^2$, then we have\n\\begin{eqnarray}\n {ds\\over d\\ln M_H} = {1\\over \\omega} - 2 s\\,,\n \\label{dsm}\n\\end{eqnarray}\nwhere $\\omega \\equiv M_H\\Omega_{ms}$ is a function of $s$ only. Eq.~(\\ref{dsm})\ncan be integrated\n\\begin{eqnarray}\n M_H(s) = M_{H,0} \\exp\\int_{s_0}^s{ds\\over\\omega^{-1}-2s}\\,,\n \\label{mh}\n\\end{eqnarray}\nwhere $M_{H,0} = M_H(s=s_0)$. Consider a Kerr black hole with initial mass $M_H$ \nand the initial spin $s = 0.998$, which is\nthe maximum value of $s$ that an astrophysical black hole can have (Thorne\n1974). As the black hole spins down to $s = 0.36$, the total amount of energy\nextracted from the black hole by the disk can be calculated with Eq.~(\\ref{mh}):\n$\\Delta E\\approx 0.15 M_H$.\nThis amount of energy will eventually be transported to infinity by the disk.\nIn a realistic case the magnetic field lines touch the disk not exactly at the \nmarginally stable orbit, the averaged angular velocity of the disk will be \nsomewhat smaller than $\\Omega_{ms}$, then the total amount of energy that \ncan be extracted\nfrom the black hole should be somewhat smaller than $0.15 M_H$.\n\nFor comparison let's calculate the amount of energy that can be extracted \nfrom a Kerr black hole by the Blandford-Znajek mechanism in the optimal case \ni.e. when the impedance matching condition is satisfied (cf. Macdonald \\&\nThorne 1982). To do so, we only need to replace $\\Omega_{ms}$ with\n$\\Omega_H/2$ in Eq.~(\\ref{mh}), since the power and torque of the black hole\nare related by $P_H = T_H\\Omega_F$ where $\\Omega_F$ is the angular velocity of \nmagnetic field lines, and in the optimal case $\\Omega_F = \\Omega_H/2$.\nThen we obtain that as the black hole spins down from $s = 0.998$ to $s = 0$ the\ntotal energy extracted from the black hole by the Blandford-Znajek mechanism\nis $\\approx 0.09 M_H$.\n\nWe find that the magnetic coupling between a black hole and a disk has a higher\nefficiency in extracting energy from the black hole than the Blandford-Znajek \nmechanism (see Fig.~\\ref{figure2}). This is because the energy extracted \nfrom the black hole by the magnetic coupling to the disk has a \nlarger ratio of energy to angular momentum than is the case for\nthe Blandford-Znajek mechanism.\n\n%\\section 4\n\\section{Conclusions}\nWhen a black hole rotates faster than the disk,\nwhich is the case if $a/M_H>0.36$ for a Kerr black hole with a \nthin Keplerian disk, then the black hole exerts a torque at the\ninner edge of the disk. The torque transfers energy\nand angular momentum from the black hole to the disk.\nThis is similar to the ``propeller'' mechanism in the case of a \nmagnetized neutron star with a disk (Illarionov \\& Sunyaev 1975). \nThe energy transfered to the disk \nis eventually radiated to infinity by the disk. This provides a mechanism for\nextracting energy from a black hole through its disk. \nFor a Kerr black hole with the initial mass $M_H$ and spin $a/M_H = 0.998$, \nthe total amount of energy that can be extracted by a thin Keplerian\ndisk is $\\approx 0.15 M_H$.\nTherefore, this is more efficient than the Blandford-Znajek mechanism\nwhich can extract only $\\approx 0.09 M_H$.\n\nWhen the black hole rotates slower than the disk, i.e. $0\\le a/M_H<0.36$,\nenergy and angular momentum are transferred from the disk to the black hole,\nand the disk accretes onto the black hole.\n\n\\acknowledgments{I am very grateful to Bohdan Paczy\\'nski for encouraging and \nstimulating discussions. This work was supported by the NASA grant NAG5-7016.}\n\n%REFERENCES\n\\begin{references}\n\n\\reference{} Begelman, M. C., Blandford, R. D., \\& Rees, M. Z. 1984, Rev.\n Mod. Phys., 56, 255\n\n\\reference{} Blandford, R. D. 1976, MNRAS, 176, 465\n\n\\reference{} Blandford, R. D. 1999, in Astrophysical Disks: An EC Summer\n School, Astronomical Society of the Pacific Conference Series,\n V. 160, ed. J. A. Sellwood \\& J. Goodman, 265\n\n\\reference{} Blandford, R. D. 2000, astro-ph/0001499\n\n\\reference{} Blandford, R. D., \\& Znajek, R. L. 1977, MNRAS, 179, 433\n\n\\reference{} Gruzinov, A. 1999, astro-ph/9908101\n\n\\reference{} Illarionov, A. F., \\& Sunyaev, R. A. 1975, A\\&A, 39, 185\n\n\\reference{} Lee, H. K., Wijers, R. A. M. J., \\& Brown, G. E. 1999, \n astro-ph/9906213\n\n\\reference{} Li, L. -X. 1999, astro-ph/9902352; to appear in\n Phys. Rev. D\n\n\\reference{} Livio, M., Ogilvie, G. I., \\& Pringle, J. E. 1999, ApJ, 512, 100\n\n\\reference{} Macdonald, D., \\& Thorne, K. S. 1982, MNRAS, 198, 345\n\n\\reference{} Novikov, I. D., \\& Thorne, K. S. 1973, in Black Holes, ed.\n C. DeWitt \\& B. S. DeWitt (NY: Gordon and Breach), 343\n\n\\reference{} Paczy\\'nski, B. 1993, in Relativistic Astrophysics and Particle\n Cosmology, ed. C. W. Akerlof \\& M. A. Srednicki, Ann. NY\n Acad. Sci., Vol. 688, 321\n\n\\reference{} Phinney, E. S. 1983, in Astrophysical Jets, ed.\n A. Ferrari \\& A. G. Pacholczyk (Dordrecht:\n\t D. Reidel Publishing Co.), 201\n\n\\reference{} Rees, M. J., Begelman, M. C., Blandford, R. D., \\& Phinney, E. S.\n 1982, Nature 295, 17\n\n\\reference{} Thorne, K. S. 1974, ApJ, 191, 507\n\n\\end{references}\n\n\\newpage\n%Figure Captions\n\\figcaption[fig1.ps]{Magnetic field lines connecting a black hole with an accretion\ndisk can transfer energy and angular momentum between them. In the figure is shown \nthe dependence of the power of the energy transfer on the spin of the black hole\nfor the model of a Kerr black hole with a thin Keplerian disk. The magnetic field lines\nare assumed to touch the disk close to the marginally stable orbit. The vertical \naxis shows the power $P_{HD}$ in unit of $P_0$, where $P_0$ is the value\nof $-P_{HD}$ for the Schwarzschild case. The horizontal axis shows $a/M_H$, where\n$M_H$ is the mass of the black hole, $a M_H$ is the angular momentum of the black\nhole. If $P_{HD}>0$, which is the case when $0.36<a/M_H <1$, energy and angular \nmomentum are transferred from the black hole to the disk; if $P_{HD}<0$, which \nis the case when $0\\le a/M_H <0.36$, energy and angular momentum are transferred from\nthe disk to the black hole. $P_{HD}$ peaks at $a/M_H\\approx 0.981$.\n\\label{figure1}}\n\n\\figcaption[fig2.ps]{The efficiency in extracting energy from a Kerr black hole\nas the black hole is spun down. The efficiency is defined by $\\eta = \\Delta E/M_H$,\nwhere $M_H$ is the mass of the black hole at its initial state with $a/M_H = 0.998$\n(the maximum value of $a/M_H$ that an astrophysical black hole can have). \n(So the left ends\nof the curves are at $a/M_H = 0.998$, not $a/M_H = 1$. Note that in the figure \n$a/M_H$ decreases from left to right.) The solid curve represents\nthe efficiency in extracting energy from a Kerr black hole through a thin Keplerian \ndisk, which ends at $a/M_H = 0.36$ since then the transfer of energy and angular \nmomentum from the black hole to the disk stops. With this mechanism, up to $\\approx \n15\\%$ of the initial mass of the black hole can be extracted. The dashed curve represents\nthe efficiency of the Blandford-Znajek mechanism, which ends at $a/M_H = 0$. With the\nBlandford-Znajek mechanism, up to $\\approx \n9\\%$ of the initial mass of the black hole can be extracted. \n\\label{figure2}}\n\n\n\\end{document}\n\n"
}
] |
[] |
astro-ph0002005
|
A Comparison of Ultraviolet, Optical, and X-Ray Imagery of Selected Fields in the Cygnus Loop
|
[
{
"author": "\\sc Charles W.\\ Danforth\\altaffilmark{1}"
},
{
"author": "Robert H.\\ Cornett\\altaffilmark{2}"
},
{
"author": "N. A. Levenson\\altaffilmark{1}"
},
{
"author": "William P.\\ Blair\\altaffilmark{1}"
},
{
"author": "Theodore P.\\ Stecher\\altaffilmark{3}"
}
] |
During the Astro-1 and Astro-2 Space Shuttle missions in 1990 and 1995, far ultraviolet (FUV) images of five 40\arcmin\ diameter fields around the rim of the Cygnus Loop supernova remnant were observed with the Ultraviolet Imaging Telescope (UIT). These fields sampled a broad range of conditions including both radiative and nonradiative shocks in various geometries and physical scales. In these shocks, the UIT B5 band samples predominantly \ion{C}{4} $\lambda$1550 and the hydrogen two-photon recombination continuum. Smaller contributions are made by emission lines of \ion{He}{2} $\lambda$1640 and \ion{O}{3}] $\lambda$1665. We present these new FUV images and compare them with optical \Ha\ and [\ion{O}{3}], and ROSAT HRI X-ray images. Comparing the UIT images with those from the other bands provides new insights into the spatial variations and locations of these different types of emission. By comparing against shock model calculations and published FUV spectroscopy at select locations, we surmise that resonance scattering in the strong FUV permitted lines is widespread in the Cygnus Loop, especially in the bright optical filaments typically selected for observation in most previous studies.
|
[
{
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\t\\@psdo\\@psfiga:=#1\\do\n\t\t {\\expandafter\\@setparms\\@psfiga,}}\n%\n% Compute bb height and width\n%\n\\newif\\ifno@bb\n\\newif\\ifnot@eof\n\\newread\\ps@stream\n\\def\\bb@missing{\n\t\\if@verbose{\n\t\t\\typeout{psfig: searching \\@p@sfile \\space for bounding box}\n\t}\\fi\n\t\\openin\\ps@stream=\\@p@sfile\n\t\\no@bbtrue\n\t\\not@eoftrue\n\t\\catcode`\\%=12\n\t\\loop\n\t\t\\read\\ps@stream to \\line@in\n\t\t\\global\\toks200=\\expandafter{\\line@in}\n\t\t\\ifeof\\ps@stream \\not@eoffalse \\fi\n\t\t%\\typeout{ looking at :: \\the\\toks200 }\n\t\t\\@bbtest{\\toks200}\n\t\t\\if@bbmatch\\not@eoffalse\\expandafter\\bb@cull\\the\\toks200\\fi\n\t\\ifnot@eof \\repeat\n\t\\catcode`\\%=14\n}\t\n\\catcode`\\%=12\n\\newif\\if@bbmatch\n\\def\\@bbtest#1{\\expandafter\\@a@\\the#1%%BoundingBox:\\@bbtest\\@a@}\n\\long\\def\\@a@#1%%BoundingBox:#2#3\\@a@{\\ifx\\@bbtest#2\\@bbmatchfalse\\else\\@bbmatchtrue\\fi}\n\\long\\def\\bb@cull#1 #2 #3 #4 #5 {\n\t\\dimen100=#2 bp\\edef\\@p@sbbllx{\\number\\dimen100}\n\t\\dimen100=#3 bp\\edef\\@p@sbblly{\\number\\dimen100}\n\t\\dimen100=#4 bp\\edef\\@p@sbburx{\\number\\dimen100}\n\t\\dimen100=#5 bp\\edef\\@p@sbbury{\\number\\dimen100}\n\t\\no@bbfalse\n}\n\\catcode`\\%=14\n%\n\\def\\compute@bb{\n\t\t\\no@bbfalse\n\t\t\\if@bbllx \\else \\no@bbtrue \\fi\n\t\t\\if@bblly \\else \\no@bbtrue \\fi\n\t\t\\if@bburx \\else \\no@bbtrue \\fi\n\t\t\\if@bbury \\else \\no@bbtrue \\fi\n\t\t\\ifno@bb \\bb@missing \\fi\n\t\t\\ifno@bb \\typeout{FATAL ERROR: no bb supplied or found}\n\t\t\t\\no-bb-error\n\t\t\\fi\n\t\t%\n\t\t\\count203=\\@p@sbburx\n\t\t\\count204=\\@p@sbbury\n\t\t\\advance\\count203 by -\\@p@sbbllx\n\t\t\\advance\\count204 by -\\@p@sbblly\n\t\t\\edef\\@bbw{\\number\\count203}\n\t\t\\edef\\@bbh{\\number\\count204}\n\t\t%\\typeout{ bbh = \\@bbh, bbw = \\@bbw }\n}\n%\n% \\in@hundreds performs #1 * (#2 / #3) correct to the hundreds,\n%\tthen leaves the result in @result\n%\n\\def\\in@hundreds#1#2#3{\\count240=#2 \\count241=#3\n\t\t \\count100=\\count240\t% 100 is first digit #2/#3\n\t\t \\divide\\count100 by \\count241\n\t\t \\count101=\\count100\n\t\t \\multiply\\count101 by \\count241\n\t\t \\advance\\count240 by -\\count101\n\t\t \\multiply\\count240 by 10\n\t\t \\count101=\\count240\t%101 is second digit of #2/#3\n\t\t \\divide\\count101 by \\count241\n\t\t \\count102=\\count101\n\t\t \\multiply\\count102 by \\count241\n\t\t \\advance\\count240 by -\\count102\n\t\t \\multiply\\count240 by 10\n\t\t \\count102=\\count240\t% 102 is the third digit\n\t\t \\divide\\count102 by \\count241\n\t\t \\count200=#1\\count205=0\n\t\t \\count201=\\count200\n\t\t\t\\multiply\\count201 by \\count100\n\t\t \t\\advance\\count205 by \\count201\n\t\t \\count201=\\count200\n\t\t\t\\divide\\count201 by 10\n\t\t\t\\multiply\\count201 by \\count101\n\t\t\t\\advance\\count205 by \\count201\n\t\t\t%\n\t\t \\count201=\\count200\n\t\t\t\\divide\\count201 by 100\n\t\t\t\\multiply\\count201 by \\count102\n\t\t\t\\advance\\count205 by \\count201\n\t\t\t%\n\t\t \\edef\\@result{\\number\\count205}\n}\n\\def\\compute@wfromh{\n\t\t% computing : width = height * (bbw / bbh)\n\t\t\\in@hundreds{\\@p@sheight}{\\@bbw}{\\@bbh}\n\t\t%\\typeout{ \\@p@sheight * \\@bbw / \\@bbh, = \\@result }\n\t\t\\edef\\@p@swidth{\\@result}\n\t\t%\\typeout{w from h: width is \\@p@swidth}\n}\n\\def\\compute@hfromw{\n\t\t% computing : height = width * (bbh / bbw)\n\t\t\\in@hundreds{\\@p@swidth}{\\@bbh}{\\@bbw}\n\t\t%\\typeout{ \\@p@swidth * \\@bbh / \\@bbw = \\@result }\n\t\t\\edef\\@p@sheight{\\@result}\n\t\t%\\typeout{h from w : height is \\@p@sheight}\n}\n%% yves\n\\def\\compute@wfroms{\n\t\t%\\typeout{computewfroms: scale is \\@p@sscale}\t\n\t\t% computing : width = scale * (bbw / 100)\n\t\t\\in@hundreds{\\@p@sscale}{\\@bbw}{100}\n\t\t%\\typeout{ \\@p@sscale * \\@bbw / 100, = \\@result }\n\t\t\\edef\\@p@swidth{\\@result}\n\t\t%\\typeout{w from s: width is \\@p@swidth}\n}\n\\def\\compute@hfroms{\n\t\t%\\typeout{computehfroms: scale is \\@p@sscale}\t\n\t\t% computing : height = scale * (bbh / 100)\n\t\t\\in@hundreds{\\@p@sscale}{\\@bbh}{100}\n\t\t%\\typeout{ \\@p@sscale * \\@bbh / 100 = \\@result }\n\t\t\\edef\\@p@sheight{\\@result}\n\t\t%\\typeout{h from s : height is \\@p@sheight}\n}\n\\def\\compute@handw{\n\t\t\\if@scale\n%\t\t\t\\edef\\@p@sheight{\\@bbh}\n%\t\t\t\\edef\\@p@swidth{\\@bbw}\n\t\t\t\\compute@wfroms\n\t\t\t\\compute@hfroms\n\t\t\\else\n\t\t\t\\if@height \n\t\t\t\t\\if@width\n\t\t\t\t\\else\n\t\t\t\t\t\\compute@wfromh\n\t\t\t\t\\fi\t\n\t\t\t\\else \n\t\t\t\t\\if@width\n\t\t\t\t\t\\compute@hfromw\n\t\t\t\t\\else\n\t\t\t\t\t\\edef\\@p@sheight{\\@bbh}\n\t\t\t\t\t\\edef\\@p@swidth{\\@bbw}\n\t\t\t\t\\fi\n\t\t\t\\fi\n\t\t\\fi\n}\n%% finish\n\\def\\compute@resv{\n\t\t\\if@rheight \\else \\edef\\@p@srheight{\\@p@sheight} \\fi\n\t\t\\if@rwidth \\else \\edef\\@p@srwidth{\\@p@swidth} \\fi\n}\n%\t\t\n% Compute any missing values\n\\def\\compute@sizes{\n\t\\compute@bb\n\t\\compute@handw\n\t\\compute@resv\n}\n%\n% \\psfig\n% usage : \\psfig{file=, height=, width=, bbllx=, bblly=, bburx=, bbury=,\n%\t\t\trheight=, rwidth=, clip=}\n%\n% \"clip=\" is a switch and takes no value, but the `=' must be present.\n\\def\\psfig#1{\\vbox {\n\t% do a zero width hard space so that a single\n\t% \\psfig in a centering enviornment will behave nicely\n\t%{\\setbox0=\\hbox{\\ }\\ \\hskip-\\wd0}\n\t%\n\t\\ps@init@parms\n\t\\parse@ps@parms{#1}\n\t\\compute@sizes\n\t%\n\t\\ifnum\\@p@scost<\\@psdraft{\n\t\t\\if@verbose{\n\t\t\t\\typeout{psfig: including \\@p@sfile \\space }\n\t\t}\\fi\n\t\t%\n\t\t\\special{ps::[begin] \t\\@p@swidth \\space \\@p@sheight \\space\n\t\t\t\t\\@p@sbbllx \\space \\@p@sbblly \\space\n\t\t\t\t\\@p@sbburx \\space \\@p@sbbury \\space\n\t\t\t\tstartTexFig \\space }\n\t\t\\if@clip{\n\t\t\t\\if@verbose{\n\t\t\t\t\\typeout{(clip)}\n\t\t\t}\\fi\n\t\t\t\\special{ps:: doclip \\space }\n\t\t}\\fi\n\t\t\\if@prologfile\n\t\t \\special{ps: plotfile \\@prologfileval \\space } \\fi\n\t\t\\special{ps: plotfile \\@p@sfile \\space }\n\t\t\\if@postlogfile\n\t\t \\special{ps: plotfile \\@postlogfileval \\space } \\fi\n\t\t\\special{ps::[end] endTexFig \\space }\n\t\t% Create the vbox to reserve the space for the figure\n\t\t\\vbox to \\@p@srheight true sp{\n\t\t\t\\hbox to \\@p@srwidth true sp{\n\t\t\t\t\\hss\n\t\t\t}\n\t\t\\vss\n\t\t}\n\t}\\else{\n\t\t% draft figure, just reserve the space and print the\n\t\t% path name.\n\t\t\\vbox to \\@p@srheight true sp{\n\t\t\\vss\n\t\t\t\\hbox to \\@p@srwidth true sp{\n\t\t\t\t\\hss\n\t\t\t\t\\if@verbose{\n\t\t\t\t\t\\@p@sfile\n\t\t\t\t}\\fi\n\t\t\t\t\\hss\n\t\t\t}\n\t\t\\vss\n\t\t}\n\t}\\fi\n}}\n\\def\\psglobal{\\typeout{psfig: PSGLOBAL is OBSOLETE; use psprint -m instead}}\n\\catcode`\\@=12\\relax\n\n\n"
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"name": "uit_pp.tex",
"string": "% Full graphical preprint version\n% latest text modification 1/21/00\n% latest format modification 1/28/00\n% chopped down size for Astro-ph archiving 1-31-00\n\n%%%*****************************************************************\n\\newcommand{\\Ha}{H$\\alpha$}\n\\newcommand{\\vel}{$\\rm km\\ s^{-1}$}\n\\newcommand{\\url}{http://www.pha.jhu.edu/$\\sim$danforth/uit/}\n\n\\documentstyle[11pt,aaspp4]{article}\n%\\documentstyle[12pt,aasms4]{article}\n\n%% EDITORIAL PERSONNEL WILL USE THE FIVE LINES BELOW.\n%% NO NEED FOR AUTHORS TO BOTHER WITH IT.\n\n%\\received{ }\n%\\accepted{ }\n%\\journalid{ }{ }\n%\\articleid{ }{ }\n%\\slugcomment{ }\n\n\\input{psfig}\n\\begin{document}\n\n\\title{A Comparison of Ultraviolet, Optical, and X-Ray Imagery of Selected\nFields in the Cygnus Loop}\n\n\\author{\\sc Charles W.\\ Danforth\\altaffilmark{1}, Robert H.\\\nCornett\\altaffilmark{2}, N. A. Levenson\\altaffilmark{1}, William P.\\\nBlair\\altaffilmark{1}, Theodore P.\\ Stecher\\altaffilmark{3}}\n\n\\altaffiltext{1}{Department of Physics and Astronomy, The Johns Hopkins\nUniversity, 3400 N. Charles Street, Baltimore, MD 21218; [email protected],\[email protected], [email protected]}\n\n\\altaffiltext{2}{Raytheon ITSS, 4400 Forbes Blvd., Lanham, MD 20706;\[email protected]}\n\n\\altaffiltext{3}{Laboratory for Astronomy and Solar Physics, NASA/GSFC, Code\n681, Greenbelt, MD 20771; [email protected]}\n\n\\begin{center}{Accepted for Publication January 27, 2000}\\end{center}\n\n\\begin{abstract}\nDuring the Astro-1 and Astro-2 Space Shuttle missions in 1990 and 1995, far\nultraviolet (FUV) images of five 40\\arcmin\\ diameter fields around the rim of\nthe Cygnus Loop supernova remnant were observed with the Ultraviolet Imaging\nTelescope (UIT). These fields sampled a broad range of conditions including\nboth radiative and nonradiative shocks in various geometries and physical\nscales. In these shocks, the UIT B5 band samples predominantly \\ion{C}{4}\n$\\lambda$1550 and the hydrogen two-photon recombination continuum. Smaller\ncontributions are made by emission lines of \\ion{He}{2} $\\lambda$1640 and\n\\ion{O}{3}] $\\lambda$1665. We present these new FUV images and compare them\nwith optical \\Ha\\ and [\\ion{O}{3}], and ROSAT HRI X-ray images. Comparing the\nUIT images with those from the other bands provides new insights into the\nspatial variations and locations of these different types of emission. By\ncomparing against shock model calculations and published FUV spectroscopy at\nselect locations, we surmise that resonance scattering in the strong FUV\npermitted lines is widespread in the Cygnus Loop, especially in the bright\noptical filaments typically selected for observation in most previous studies.\n\n\\end{abstract}\n\n\\keywords{ISM: nebulae --- ISM: supernova remnants --- ISM: shock waves ---\nultraviolet: imaging}\n\n\\section{Introduction}\nBecause of its large angular size and wide range of shock conditions, the\nCygnus Loop is one of the best laboratories for studying the environment and\nphysics of middle-aged supernova remnants (SNR). It covers a huge expanse in\nthe sky (2.8$\\times$3.5$\\rm ^o$) corresponding to 21.5$\\times$27 pc, at a newly\ndetermined distance of 440 pc (\\cite{Blair99}). The currently accepted view\nfor the Cygnus Loop is that it represents an explosion in a cavity produced by\na fairly massive precursor star (cf. \\cite{Levenson98}). The SN shock has\nbeen traveling relatively unimpeded for roughly ten parsecs and has only\nrecently begun reaching the denser cavity walls. The size of the cavity\nimplicates a precursor star of type early B. The interaction of the shock with\nthe complex edges of the cavity wall is responsible for the complicated mixture\nof optical and X-ray emission seen in superposition, and a dazzling variety of\noptical filament morphologies.\n\nPortions of the SN blast wave propagating through the fairly rarefied atomic\nshell ($<$1 cm$^{-3}$), show faint filaments with hydrogen\nBalmer-line-dominated optical spectra. These filaments represent the position\nof the primary blast wave and are often termed nonradiative shocks (because\nradiative losses are unimportant to the dynamics of the shock itself). Ambient\ngas is swept up and progressively ionized, emitting \\ion{He}{2}, \\ion{C}{4},\n\\ion{N}{5}, and \\ion{O}{6} lines in the FUV (Figure~1, bottom spectrum)\n(\\cite{Hester94}, \\cite{Raymond83}). Balmer-dominated emission arises from the\nfraction ($\\sim$0.3) of neutral hydrogen swept up by the shock that stands some\nchance of being excited and recombining before it is ionized in the post-shock\nflow (\\cite{Chevalier78}; \\cite{CKR80}).\n\nThe Balmer emission is accompanied by hydrogen two-photon events which produce\na broad continuum above 1216\\AA\\ peaking at $\\sim$1420\\AA\\\n(\\cite{Nussbaumer84}). For recombination and for high temperature shocks, the\nratio of two-photon emission to Balmer is nearly constant ($\\sim$8:1). In\nslow shocks ($\\sim$40\\vel) in neutral gas, the ratio can be enhanced\nconsiderably (\\cite{Dopita82}).\n\nBalmer-dominated filaments are very smooth and WFPC2 observations by Blair et\nal. (1999) show that they are exceedingly thin as well---less than one WFC\npixel across when seen edge-on, or $<6 \\times 10^{14}$ cm at our assumed\ndistance, in keeping with theoretical predictions (cf. \\cite{Raymond83}).\nPostshock temperatures reach millions of degrees and the hot material emits\ncopious soft X-rays. The density is low, however, and cooling is very\ninefficient. With time, as the shock continues to sweep up material, these\nfilaments will be able to start cooling more effectively and will evolve to\nbecome radiative filaments.\n\nThe bright optical filaments in the Cygnus Loop represent radiative shocks in\nmuch denser material, such as might be expected in the denser portions of the\ncavity wall. These shocks are said to be radiative (that is, energy losses\nfrom radiation are significant); they have more highly developed cooling and/or\nrecombination zones. The shocked material emits in the lines of a broad range\nof hot, intermediate, and low temperature ions, depending on the effective\n`age' of the shock at a given location and the local physical conditions. For\ninstance, a relatively recent encounter between the shock and a density\nenhancement (or similarly, a shock that has swept up a fairly low total column\nof material) may show very strong [O~III] $\\lambda 5007$ compared with \\Ha.\nThis would indicate that the coolest part of the flow, the recombination zone\nwhere the Balmer lines become strong, has not yet formed. Such shocks are said\nto be `incomplete' as the shocked material remains hot and does not yet emit in\nthe lower ionization lines.\n\nIn contrast, radiative filaments with the full range of ionization (including\nthe low ionization lines) are well approximated by full, steady-flow shock\nmodel calculations, such as those of \\cite{Raymond79}, \\cite{Dopita84}, and\nHartigan, Raymond, \\& Hartmann (1987; hereafter \\cite{HRH}). Morphologically,\nradiative complete filaments lack the smooth grace of nonradiative filaments\nor even radiative incomplete filaments in some cases (cf. \\cite{Fesen82}). The\nmore irregular appearance of these filaments is due partly to inhomogeneities\nin the shocked clouds themselves, partly to turbulence and/or thermal\ninstabilities that set in during cooling (cf. \\cite{Innes92} and references\ntherein), and partly to several clouds appearing along single lines of sight.\nOften the emission at a given filament position cannot be characterized by a\nsingle shock velocity.\n\nMuch of the above understanding of shock types and evolutionary stages has been\npredicated on UV/optical studies of the Cygnus Loop itself. The Cygnus Loop is\n a veritable laboratory for such studies because of its relative proximity,\nlarge angular extent and low foreground extinction (E[B $-$ V] = 0.08;\n\\cite{Fesen82}), and thus its accessibility across the electromagnetic\nspectrum. However, because of the range of shock interactions and shock types,\ncoupled with the significant complication of projection effects near the limb\nof the SNR, great care must be taken in order to obtain a full understanding of\nwhat is happening at any given position in the nebular structure.\n\nAlthough FUV spectra are available at a number of individual filament locations\nfrom years of observations with IUE and the shuttle-borne Hopkins Ultraviolet\nTelescope (HUT), the perspective obtainable from FUV imaging has been largely\nlacking. The Ultraviolet Imaging Telescope (UIT) was flown as part of the\nAstro-1 Space Shuttle mission in 1990 was used to observe a field in the Cygnus\nLoop through both mid-UV and far-UV (FUV) filters (\\cite{Cornett92}). In this\npaper, we report on additional FUV observations with UIT obtained during the\nAstro-2 shuttle mission in 1995. In addition to the field imaged during\nAstro-1, UIT observed four different regions around the periphery of the Cygnus\nLoop with a resolution comparable to existing optical and X-ray observations.\nThese fields sample the full range of physical and shock conditions and\nevolutionary stages in the SNR. We combine these data with existing\nground-based optical images and ROSAT HRI X-ray data to obtain new insights\ninto this prototypical SNR and its interaction with its surroundings.\n\nIn \\S2 we present the observations obtained with the UIT and review the\ncomparison data sets. In \\S3 we discuss the spectral content of the UIT filter\nused in the observations. In \\S4 we discuss examples of the various kinds of\nshocks as seen in the UIT fields, and summarize our conclusions in \\S5.\n\n\\section{UIT Observations and Comparison Data}\n\nUIT has flown twice on the Space Shuttle as part of the Astro-1 and Astro-2\nprograms (1990 December 2-10 and 1995 March 2-18). Together with the Hopkins\nUltraviolet Telescope (HUT) and the Wisconsin Ultraviolet Photo-Polarimeter\nExperiment, UIT explored selected UV targets. An f/9 Ritchey-Chretien\ntelescope with a 38 cm aperture and image intensifier systems produced images\nof circular 40\\arcmin\\ fields of view with $\\sim$3\\arcsec\\ resolution at field\ncenter (depending on pointing stability). Images were recorded on 70mm Eastman\nKodak IIa-O film which was developed and digitized at NASA/GSFC and processed\ninto uniform data products. Technical details on the hardware and data\nprocessing can be found in \\cite{Stecher92} and Stecher et al. (1997).\n\n%\\begin{table}\n\\noindent Table~1: UIT B5 Filter Observations in the Cygnus Loop\\\\\n\\begin{tabular}{lllll}\nPosition & RA(J2000)& Dec(J2000)& exposure (sec) & Figure\\\\\n\\hline\nW cloud & 20:45:38 & +31:06:33 & 1010 & 3\\\\\nNE nonrad & 20:54:39 & +32:17:29 & 2041 & 4\\\\\nNE cloud & 20:56:16 & +31:44:34 & 500 & 5\\\\\nXA region$\\rm^a$& 20:57:35 & +31:07:28 & 1280 & 6\\\\\nXA region & 20:57:04 & +31:07:45 & 1151 & 6\\\\\nXA region & 20:57:22 & +31:04:02 & 1516 & 6\\\\\nXA region & 20:57:24 & +31:03:51 & 1274 & 6\\\\\nSE cloud & 20:56:05 & +30:44:01 & 2180 & 7\\\\\n\\hline\n\\end{tabular}\\\\\n$\\rm^a$ Astro-1 image (cf. \\cite{Cornett92})\n\nAstro/UIT images are among the few examples of FUV images of SNRs, and UIT's B5\nbandpass ($\\sim1450$\\AA\\ to $\\sim1800$\\AA) encompasses severally generally\nhigh-excitation and heretofore unmapped lines that are often present in SNR\nshocks (Figure~1). UIT's two Astro flights have produced eight FUV images of\nfive different Cygnus Loop fields. Table~1 lists the observation parameters\nand field locations, which are indicated in Figure~2. We will refer to these\nfields by the names listed in Table~1. Since all four exposures of the XA\nregion (named by \\cite{HesterCox86}) are reasonably deep, we constructed a\nmosaic of the field using the IRAF\\footnotemark\\ IMCOMBINE task, resulting in\nsignificantly improved signal-to-noise in the overlapped region of the combined\nimage. In panel c of Figures 3 through 7, we show the five reduced UIT images\nas observed in the B5 filter bandpass. \\footnotetext{IRAF is distributed by\nthe National Optical Astronomy Observatories, which is operated by the\nAssociation of Universities for Research in Astronomy, Inc.\\ (AURA) under\ncooperative agreement with the National Science Foundation.}\n\nUIT images with long exposure times suffer from an instrumental malady dubbed\n``measles'' by the UIT team (\\cite{Stecher97}). Measles manifest themselves as\nfixed-pattern noise spikes in images with a large sunlight flux, such as long\ndaylight exposures or images of red, very bright sources (e.g. planets or the\nMoon). This effect is probably produced by visible light passing through\npinholes in either the output phosphor of the first stage or the bialkali\nphotocathode of the second stage of the UIT FUV image tube. The Cygnus Loop\nwas a daytime object for both the Astro-1 and Astro-2 flights, but the\nphenomenon is visible only in some of the longer exposures. Most dramatically,\nmeasles are seen in the northeast cloud nonradiative image (Figure~4) as a\ndarkening in the northwest corner; the individual ``measles'' are spread into a\nbackground by the binning used to produce these images. Various approaches to\nremoving the appearance of measles were attempted but none of them have yielded\nsatisfactory results. In practice, the measles, here arising from daylight sky\ncontamination, affect our analysis only by adding to the background level, so\nthe original images are presented here, ``measles'' and all.\n\nFor comparison with our FUV images, we show narrow-band optical images in\n[\\ion{O}{3}] $\\lambda$5007 and \\Ha+[\\ion{N}{2}] (which for simplicity we refer\nto as \\Ha) obtained with the Prime Focus Corrector on the 0.8 m telescope at\nMcDonald Observatory (cf. \\cite{Levenson98}). These images, shown in panels~a\nand b of Figures~3~--~7, are aligned and placed on a common scale of 5\\arcsec\\\nper pixel, which is similar to the FUV resolution of 3\\arcsec. The optical\nimages have each been processed with a 3-pixel median filter to remove faint\nstars and stellar residuals.\n\nIn addition, we show the soft X-ray (0.1--2.4 keV) emission for each field, as\nobserved with the {\\it ROSAT} High Resolution Imager (HRI) (from\n\\cite{Levenson97}). The resolution of the HRI imager is 6\\arcsec\\ on axis\ndegrading to 30\\arcsec\\ at the edge of each field. As with the optical data,\nthe X-ray images are aligned on a 5\\arcsec\\ per pixel scale. The HRI images\nhave additionally been smoothed with a 3-pixel FWHM gaussian and are shown in\npanel d of Figures~3~--~7. All images in Figures~3~--~7 are displayed on a\nlogarithmic scale.\n\nFigure~8 shows three-color composite images using \\Ha\\ as red, B5 as green, and\nthe {\\it ROSAT} HRI as blue. The color levels have been adjusted for visual\nappearance, to best show the relative spatial relationships of the different\nemissions. (The color composite for the Northeast nonradiative region is not\nshown, since little new information is gained above Figure~4 and because of the\nadverse effect of the measles.) This will be discussed in more detail below.\n\n\\section{Spectral Content of UIT Images}\n\nFigure~1 shows the UIT B5 filter profile superimposed on spectra of typical\nradiative and nonradiative filaments, as observed by HUT. Unlike typical SNR\nnarrow band images in the optical, the B5 filter is relatively broad and does\nnot isolate a single spectral line, but rather encompasses several strong,\nmoderately high ionization lines that are variable from filament to filament.\nCornett et al. (1992) point out that \\ion{C}{4} should dominate emission in\nthis bandpass since it is a strong line centered near the filter's peak\nthroughput, and since shock models predict this result for a range of important\nvelocities (cf. Figure~10 and accompanying discussion). Here we look at this\nmore closely over a larger range of shock velocities, and in particular also\ndiscuss the potential complicating effects of hydrogen two-photon recombination\ncontinuum emission, shock completeness, and resonance line scattering.\n\nEmpirical comparisons of IUE and HUT emission line observations can be used to\nquantify at what level the C~IV emission is expected to dominate the line\nemission detected through the B5 filter. For instance, in the highly radiative\nXA region (see Figure~7) we have compared a large number of FUV spectra both\non and adjacent to bright optical filaments against the throughput curve of B5\n(Danforth, Blair, \\& Raymond 2000; henceforth \\cite{DBR}). This comparison\nshows that on average the various lines contribute as follows:\n\\ion{C}{4} $\\lambda$1550, 42\\%;\n\\ion{O}{3}] $\\lambda$1665, 27\\%;\n\\ion{He}{2} $\\lambda$1640, 17\\%;\n\\ion{N}{4} $\\lambda$1486, 8\\%;\nand 6\\%\nfrom fainter emission lines. Using the HUT observation of \\cite{Long92}, we\nestimate for nonradiative shocks the B5 contributions are more like\n\\ion{C}{4} (60\\%),\n\\ion{He}{2} (28\\%),\nand all other species 12\\%.\nThese percentages are only approximate, of course, and will vary with shock\nvelocity, geometry and a host of other conditions, but they serve to highlight\nthe fact that, while \\ion{C}{4} is the strongest contributor to the line\nemission, it is not the only contributor.\n\nIn addition, while it is not obvious at the scale of Figure~1, a low level\ncontinuum is often seen in IUE and HUT spectra of Cygnus Loop filaments,\nespecially where optical \\Ha\\ emission is present and strong. This continuum\narises due to the hydrogen two-photon process (cf. \\cite{Osterbrock89}).\n\\cite{Benvenuti80} note that SNR shocks cause two-photon emission from hydrogen\nvia both collisional excitation and recombination into the 2$^2$S$_{1/2}$\nstate. The two-photon spectrum arises from a probability distribution of\nphotons that is symmetric about 1/2 the energy of Ly$\\alpha$ (corresponding to\n2431\\AA), resulting in a shallow spectral peak near 1420\\AA\\ and extending from\n1216\\AA\\ towards longer wavelengths, throughout the UV and optical region. The\nexpected (integrated) strength of this component is about 8 $\\times$ the \\Ha\\\nflux but is spread over thousands of Angstroms. However, the wide bandpass of\nthe B5 filter detects $\\sim$15\\%\nof the total two-photon flux available, enough to compete with line emission in\nthe bandpass. Further complicating the question, two-photon emission can also\nbe highly variable from filament to filament.\n\nBy using signatures from the images and spectra at other wavelengths, we can\ninterpret, at least qualitatively, what is being seen in the UIT images. For\ninstance, Figure~4 shows the NE rim of the SNR. The faint, smooth \\Ha\\\nfilament running along the edge of the X-ray emission is clearly a nonradiative\nfilament. The faint emission seen in the B5 image traces these faint Balmer\nfilaments well, and at this position, does not correlate particularly well with\nthe clumpy [\\ion{O}{3}] emission seen near the middle of the field. This\nimplies a relatively strong contribution from two-photon continuum, although as\nshown in the bottom spectrum of Figure~1, \\ion{C}{4} and \\ion{He}{2} are also\npresent in the filaments at some level.\n\nAs discussed earlier, in radiative filaments, higher ionization lines such as\n\\ion{O}{6} $\\lambda$1035, \\ion{N}{5} $\\lambda$1240, \\ion{C}{4} $\\lambda$1550,\nand [\\ion{O}{3}] $\\lambda$5007 become strong first, followed by lower\nionization lines like [\\ion{S}{2}] $\\lambda$6725, [\\ion{O}{1}] $\\lambda$6300,\nand the hydrogen Balmer lines. Hence, in filaments that show high optical\n[\\ion{O}{3}] to \\Ha\\ ratios, and are thus incomplete shocks, the B5 content\nprimarily arises from \\ion{C}{4} and other line emission. In older, more\ncomplete shocks where the optical [\\ion{O}{3}] to \\Ha\\ ratios are close to\nthose expected from steady flow shock models, two-photon emission again should\ncompete with the line emission and the B5 flux should arise from both sources.\nIt is difficult to assess these competing effects from Figures 3 -- 7 since the\nrelative intensities of the two optical images are not always obvious, but much\nof the variation in coloration in Figure~8 for bright radiative filaments is\ndue to the variation in relative amounts of line emission and two-photon\ncontinuum contributions to the B5 image.\n\nIn Figure~9, we show the XA field as seen with UIT (panel a) and ratio maps of\nthe UIT image against the aligned optical \\Ha\\ and [\\ion{O}{3}] images. Since\nthe ionization energies of \\ion{C}{4} (64.5 eV) and \\ion{O}{3} (54.9 eV) are\nsimilar (and to the extent that the B5 image contains a substantial component\nof \\ion{C}{4} emission), we would expect a ratio of B5 to \\Ha\\ to show evidence\nfor the transition from incomplete to complete shock filaments. Such a ratio\nmap is shown in panel b of Figure~9, and a systematic pattern is indeed seen.\nThe white filaments, indicative of a relatively low value of the ratio (and\nhence relatively strong \\Ha\\ filaments) tend to lie systematically to the\nright. These filaments tend to be closer to the center of the SNR, and hence\nshould have had more time (on average) to cool and recombine. Of course, there\nis significant evidence for projection effects in this complicated field as\nwell. Indeed, one interpretation of Figure~9b is that we are separating some\nof these projection effects, and are seeing two separate `systems' of filaments\nthat are at differing stages of completeness.\n\nAnother way of assessing the expected contributions of line emission and\ntwo-photon emission to the B5 flux is by comparing to shock model calculations.\nWe use the equilibrium preionization ``E'' series shock models of Hartigan,\nRaymond, \\& Hartmann (1987, \\cite{HRH}) to investigate variations in spectral\ncontributions to the UIT images as a function of shock velocity. Figure~10\nshows how various spectral components are predicted to change in relative\nintensity as shock velocity increases for this set of planar, complete, steady\nflow shock models. As expected, the key contributors to the B5 bandpass are\nindeed \\ion{C}{4} and two-photon continuum, although between $\\sim$100 --\n200~\\vel\\ these models indicate \\ion{C}{4} should dominate.\n\nThis is quite at odds with `ground truth', as supplied by careful comparisons\nat the specific locations of IUE and HUT spectra within the UIT fields of\nview. We note that the two-photon flux per \\AA\\ in HUT and IUE spectra is low\nand thus difficult to measure accurately since background levels are poorly\nknown. Even so, it is quite clear from comparisons such as those of\n\\cite{Benvenuti80} and Raymond et al. (1988) that nowhere do we see \\ion{C}{4}\ndominate at the level implied by Figure~10. (Indeed such studies indicate that\ntwo-photon should dominate! As will be discussed more thoroughly in \\S4,\n\\cite{Benvenuti80} and others give two-photon fluxes which overwhelm \\ion{C}{4}\nin the B5 band by a factor of 5-10. Interestingly, consideration of\nincompleteness effects only serves to exacerbate this discrepancy since the\nexpected two-photon emission should be weaker or absent. Something else is\ngoing on.\n\nThat `something else' is apparently resonance line scattering. It has long\nbeen suspected that the strong UV resonance lines, like \\ion{N}{5}\n$\\lambda$1240, \\ion{C}{2} $\\lambda$1335, and \\ion{C}{4} $\\lambda$1550, are\naffected by self-absorption along the line of sight, either by local gas within\nthe SNR itself or by the intervening interstellar medium. We can expect\nsignificant column depth from the cavity wall of the remnant itself. Since\nfilaments selected for optical/UV observation have tended to be bright, and\nsince many such filaments are edge-on sheets of gas with correspondingly high\nline of sight column densities (\\cite{Hester87}), the spectral observations are\nlikely affected in a systematic way.\n\nWhile this has been known for some years (\\cite{Raymond81}), the UIT data\npresented here indicate just how widespread resonance line scattering is in the\nCygnus Loop and how significantly the \\ion{C}{4} intensity may be reduced by\nthis effect. Figure~9c shows a ratio map of the B5 image to the [\\ion{O}{3}]\noptical image of the XA region (cf. \\cite{Cornett92}). Since [\\ion{O}{3}] is\na forbidden transition, its optically thin emission is not affected by\nresonance scattering. The ionization potentials for \\ion{C}{4} and\n[\\ion{O}{3}] are similar, so this ratio should provide some information about\nresonance scattering, if a significant fraction of the B5 image can be\nattributed to \\ion{C}{4}. Hence, this ratio image shows where resonance line\nscattering is most important, and provides information on the 3-dimensional\nstructure of regions within the SNR.\n\nThe B5 image gives the {\\em appearance} of smaller dynamic range and lower\nspatial resolution than [O~III] because we see optically thick radiation from\nonly a short distance into the filaments. The highest saturation (lowest\nratios, or light areas in Figure~9c) occurs in the cores of filaments and dense\nclouds, such as the three regions indicated in Figure~9a. The ``spur''\nfilament was studied in detail by \\cite{Raymond88} and is probably an edge-on\nsheet of gas. The region marked `B' is the turbulent, incomplete shocked cloud\nobserved with HUT during Astro-1 (\\cite{Blair91}). The XA region is also a\nshocked cloud or finger of dense gas that is likely elongated in our line of\nsight (cf. \\cite{HesterCox86}; \\cite{DBR}). What is surprising, however, is the\nextent to which the light regions in Figure~9c extend beyond the cloud cores\ninto regions of more diffuse emission. This indicates that significant\nresonance scattering is very widespread in the Cygnus Loop. The diminished\n\\ion{C}{4} flux also boosts the relative importance of two-photon emission in\nthe B5 bandpass and explains the discrepancy between numerous spectral\nobservations and the shock model predictions shown in Figure~10.\n\nUIT's B5 images are particularly useful in that they sample two important shock\nphysics regimes--the brightest radiative shocks arising in dense clouds and the\nprimary blast wave at the edge of the shell. However, it is evidently\ndifficult to predict the spectral content of B5 images alone without detailed\nknowledge of the physics of the emitting regions. Nonetheless, B5 images are\nuseful in combination with [\\ion{O}{3}]$\\lambda$5007 and \\Ha\\ images as\nempirical tools. The image combinations allow us to determine whether\n\\ion{C}{4} or two-photon dominates, in two clear-cut cases. 1) In regions\nwhere B5 images closely resemble [\\ion{O}{3}] images, the B5 filter is\ndetecting radiative shocks with velocities in the range 100-200 \\vel\\ and\ntherefore primarily \\ion{C}{4}. 2) In regions where B5 images closely resemble\n\\Ha, the B5 filter is detecting largely two-photon emission from recombination\nof hydrogen in radiative shocks or from collisional excitiation of hydrogen in\nnonradiative shocks.\n\n\\section{Discussion}\n\nEach of the fields in our study portrays a range of physical conditions and\ngeometries, and hence filament types, seen in projection in many cases. By\ncomparing the UV, X-ray and optical emissions, we can gain new insights into\nthese complexities. In this section, we discuss the spatial relationships\nbetween the hot, intermediate and cooler components seen in these images.\n\n\\subsection{The Western Cloud}\n\nIn the Western Cloud field (Figure~3) the B5 image of the bright north-south\nfilaments resembles the [\\ion{O}{3}]5007\\AA~ images very closely. The\nfilaments are clearly portions of a radiative shock viewed edge-on to our line\nof sight. The Western Cloud has been studied spectroscopically at optical\nwavelengths by \\cite{Miller74} and in the FUV by \\cite{Raymond80b}.\n\nThis region shows a case where a cloud is evidently being overrun by a shock,\nand the cloud is much larger than the scale of the shock. The cloud is\nelongated in the plane of the sky of dimensions perhaps 1$\\times$10 pc\n(\\cite{Levenson98}) and represents an interaction roughly 1000 years old\n(\\cite{Levenson96}). The main north-south radiative filament is bright in all\nwavelengths, with good detailed correlation between B5 and [\\ion{O}{3}]. \\Ha\\\nis seen to extend farther to the east, toward the center or 'behind' the shock,\nas is expected in a complete shock stratification.\n\nBright X-rays (Figure~3d) are seen to lag behind the radiative filaments by\n1-2\\arcmin\\ (0.15 to 0.3 pc). This is indicative of a reverse shock being\ndriven back into the interior material from the dense cloud. This\ndoubly-shocked material shows enhanced brightness of about a factor of 2. From\nthis, Levenson et al. (1996) derive a cloud/ambient density contrast of about\n10.\n\nAttempts to fit shock models to optical observations of the bright filament\nhave been frustrated by the large [\\ion{O}{3}]/H$\\beta$ ratio. A shock\nvelocity of 130 \\vel\\ was found by \\cite{Raymond80b} using IUE line strengths\nand assuming a slight departure from steady flow and depleted abundances in\nboth C and Si. \\cite{Raymond80b} also note that much of the hydrogen\nrecombination zone predicted by steady flow models is absent, implying that the\ninteraction is fairly young.\n\nAs seen in the \\Ha\\ image, a Balmer-dominated filament projects from the south\nof the bright radiative filament toward the northwest. \\cite{Raymond80a} find\nthat the optical spectrum of the filament contains nothing but hydrogen Balmer\nlines. High-resolution observations of the \\Ha\\ line (\\cite{Treffers81}) show\na broad component and a narrow component, corresponding to the pre- and\npost-shock conditions in the filament, with a resulting estimated shock\nvelocity of 130-170 \\vel. The filament may be a foreground or background piece\nof the blast wave not related to the radiative portion of the shock, or a\nrelated piece of blast wave that is travelling through the atomic (rather than\nmolecular) component. It is visible in both \\Ha\\ (Figure~3a) and B5\n(Figure~3c) though generally not in other bands; thus the B5 flux for this\nfilament arises primarily from the two-photon process. There is a small\nsegment of the filament visible in [\\ion{O}{3}] where the shock may be becoming\nradiative, visible in B5 as a brightening near the southern end of the\nfilament.\n\nThe X-ray luminosity behind this nonradiative filament is much fainter than\nthat observed to the east of the main radiative filament, since there is no\nreverse shock associated with the nonradiative filament to boost the brightness\n(\\cite{Hester94}). The absence of X-rays to the west of this filament confirms\nthat it represents the actual blast front. As expected, the peak X-ray flux\nlags behind the \\Ha\\ and B5 flux by roughly one arcminute (0.1 pc). This\nrepresenting the ``heating time'' of gas behind the shock.\n\nA CO cloud is seen just to the south of the Western Cloud field\n(\\cite{Scoville77}). The presence of CO clearly indicates material with\nmolecular hydrogen at densities of 300-1000 cm$^{-3}$. The nonradiative\nfilament runs closely along the T$_{antenna}$=5K contour of the CO cloud,\nindicating this shock is moving through the atomic component at this stage, but\nshowing no sign of interaction with the molecular cloud.\n\n\\subsection{Northeast Nonradiative Region}\n\nThe canonical example of nonradiative filaments in any context lies on the\nnorth and northeast rim of the Cygnus Loop. There, smooth Balmer filaments\nextend counterclockwise from the northern limb (Figure~4), and can be seen\nprominently in \\Ha\\ in Figure~5a. Small portions of this shock system have\nbeen extensively studied by Raymond et al. (1983), Blair et al. (1991),\n\\cite{Long92}, Hester, Raymond \\& Blair (1994), and most recently by Blair et\nal. (1999). The filaments are clearly visible in \\Ha\\ (Figure~4a) as well as\nB5 (Figure~4c), but invisible along most of their length in [\\ion{O}{3}]\n(Figure~4b) except for small segments. These segments represent portions of the\nshock front where a slightly higher density has allowed the shock to become\npartially radiative. The shocked, T$\\sim10^{6}$K gas emits in an\nedge-brightened band of X-rays (Figure~4d). The brightness variations in\nX-rays confirm that the nonradiative filaments are simply wrinkles in the blast\nwave presenting larger column densities to our line of sight.\n\nSpectroscopic observations of selected locations on the filaments indicate that\nthe B5 filter observes nonradiative filaments as a mixture of \\ion{C}{4} and\ntwo-photon emission. \\cite{Long92} find an intrinsic ratio of two-photon\nemission to \\ion{C}{4} of 4.3, which gives an observed ratio in B5 of 0.65.\nRaymond et al. (1983) find fluxes in the same filament which give an observed\nratio of 1.6; in a nearby filament, Hester, Raymond \\& Blair (1994) find a\nratio near 2.0. These filaments all have velocities of around 170 \\vel. It is\nlikely that much of the ISM carbon is locked up in grains in the preshock\nmedium, thus boosting the ratio.\n\nThe system of thin filaments in the NE nonradiative field extends to the south\nand is visible in \\Ha\\ ahead of the radiative Northeast Cloud (Figure~5)\ndiscussed below.\n\n\\subsection{The Northeast Cloud}\n\nThe Northeast Cloud (Figure~5) radiative filaments, south and east of the field\ndiscussed above, make up one of the brightest systems in the Cygnus Loop. The\ninteraction of the SN blast wave and the denser cavity wall is most evident at\nthis location. A complex of radiative filaments can be seen, apparently\njumbled together along our line of sight, displaying the signs of a complete\nshock undergoing radiative cooling. The X-ray edge marking the SN blast wave\nis well separated from the optical and UV filaments, implying a strongly\ndecelerated shock and cooling that has continued for some time. Stratification\nof different ionic species is evident, with [\\ion{O}{3}] in sharp filaments to\nthe east, and more diffuse \\Ha\\ behind (Figure~8b).\n\nThe Northeast Cloud extends into the southern portions of the NE nonradiative\nfield (Figure~4) as well. However, the exposure time for this FUV image is a\nfactor of four shorter than that in Figure~4c, so the nonradiative filaments\nare not detected above the background. There are a few UV-bright sections which\ncorrespond closely with bright [\\ion{O}{3}] knots. However, other equally\nbright [\\ion{O}{3}] knots in the region do not have corresponding FUV knots.\nThis may be evidence for a range of shock velocities, or it may be portions of\nthe shocks that are in transition from nonradiative to radiative conditions.\n\nUsing IUE spectra \\cite{Benvenuti80} measure the two-photon continuum for one\nof the brightest radiative positions within the NE cloud, with a resulting\nobserved two-photon/\\ion{C}{4} ratio of 5.0. Observations of other radiative\nregions both in the Cygnus Loop and in other SNRs similar in morphology and\nspectrum give ratios between 1.7 and 10 (\\cite{Raymond88}; various unpublished\ndata). Therefore, while conditions vary widely within these shocked regions,\nspectroscopy indicates that resonance scattering of \\ion{C}{4} causes us to see\n2-6 times more flux from two-photon emission than from other ions in the field.\n Yet the B5 morphology of most of the field resembles [\\ion{O}{3}] far more\nthan \\Ha, as we would expect if two-photon emission were dominant. The\napparent conflict is likely caused by the fact that most lines of sight through\nthis region undoubtedly encounter material with a broad range of physical\nconditions. Furthermore, the UIT NE cloud exposure is the shortest of our set.\n Only regions bright in both \\Ha\\ and [\\ion{O}{3}] show up in B5.\n\n\\subsection{The XA Field}\n\nThe XA field (Figure~6) is a complicated region of predominantly radiative\nfilaments, noteworthy because an extremely bright and sharp X-ray edge\ncorresponds closely to a bright knot of UV/visible emission\n(\\cite{HesterCox86}). Indeed, this region is seen to be bright in many\nwavelengths including radio (\\cite{Green90}; \\cite{Leahy97}) and infrared\n(\\cite{Arendt92}). Strong \\ion{O}{6} $\\lambda$1035 emission is seen\n(\\cite{Blair91}) as well as other high-ionization species; \\ion{N}{5},\n\\ion{C}{4}, \\ion{O}{3}] (\\cite{DBR}) and [\\ion{Ne}{5}] (\\cite{Szentgyorgyi99}).\nSee DBR for a more detailed analysis of this region.\n\nIn general, the B5 emission corresponds closely to optical [\\ion{O}{3}].\nHowever, while optical images show a high contrast between the brightest\n`cloud' regions and others in 'empty' space, B5 contrast is lower\n(\\cite{Cornett92}). This suggests contributions from a high column depth of\ndiffuse \\ion{C}{4} and/or two-photon emission. We are either looking at\ndiffuse material through the edge of a cavity wall or are seeing emission from\nface on sheets of gas. \\cite{DBR} show evidence that the bright 'cloud' in the\ncenter is not isolated and may be a density enhancement in the cavity wall or a\nfinger of denser material projecting in from the east. The entire blast wave\nin the region appears indented from the otherwise circular extent of the SNR\n(\\cite{Levenson97}) implying that the disturbance is produced by a cloud\nextended several parsecs in our line of sight. The visible structure is likely\nthe tip of a much larger cloud.\n\nLevenson et al. (1998) suggest a density enhancement in the cavity wall,\nresulting in rapid shock deceleration and accounting for the bright emission.\nIUE and HUT observations show evidence for a 150 \\vel\\ cloud shock in the dense\ncore of XA itself (the west-pointing V shape in the center of the field) and a\nfaster, incomplete shock in the more diffuse regions to the north and south\n(\\cite{DBR}). Two parallel, largely east-west filaments are seen flanking the\ncentral 'cloud'. The X-ray emission is seen to drop off dramatically south of\nthe two long radial filament systems.\n\nBlair et al. (1991) report HUT observations of a radiative but incomplete cloud\nshock directly to the north of XA marked `B' in Figure~9a. This region\nfeatures almost complete cooling with the exception of \\Ha\\ and cooler ions.\nRaymond et al. (1988) studied the Spur filament and found a completeness\ngradient along the length of it. This filament is well-defined in B5 as well\nas the optical bands.\n\nThe XA region is the one region in the Cygnus Loop where preionization is\nvisible ahead of the shock front (\\cite{Levenson98}). This preionization is\ncaused by X-ray flux from the hot, postshock gas ionizing neutral material\nacross the shock front. The emission measure is high enough in this\nphotoionized preshock gas that it is clearly visible as a diffuse patch of\nemission a few arc minutes to the east of the main XA knot in the center of the\nfield in both \\Ha\\ and B5. The B5 flux presumably arises almost entirely from\ntwo-photon emission in this case since no [\\ion{O}{3}] is seen (and hence no\nstrong UV line emission is expected).\n\nOne unique ability of the B5 filter becomes apparent in the XA region; that of\ndetecting nonradiative shocks in ionized gas. In the X-ray (Figure~6d) we see\na bulge of emission to the north and east of the brightest knot (Hester \\&\nCox's XA region proper). This bulge does not show up in either of the optical\nbands, but the perimeter is visible in the FUV at the edge of the X-ray\nemission in Figure~6c. This region has likely been ionized by X-ray flux from\nthe hot post-shock gas. A nonradiative shock is now propagating through it\nand, lacking a neutral fraction to radiate in \\Ha, is seen only in high ions\nsuch as \\ion{C}{4}. This filament is becoming more complete in its southern\nextremity (the `B' location in Figure~9a) and is emitting in [\\ion{O}{3}] as\nwell. This filament also appears to connect to the nonradiative filament seen\nin \\Ha\\ in the Northeast cloud (Figure~5a).\n\n\\subsection{The Southeast Cloud}\n\nThe Southeast Cloud (Figure~7) presents an interesting quandry. In the optical\nit appears as a small patch of radiative emission with a few associated\nnonradiative filaments. \\cite{Fesen92} hypothesize that it represents a small,\nisolated cloud at a late stage of shock interaction. Indeed, the resemblance\nto the late-stage numerical models of \\cite{Bedogni90} and \\cite{Stone92} is\nstriking.\n\nMore recent X-ray analysis (\\cite{Graham95}) suggests that the shocked portion\nof the southeast cloud is merely the tip of a much larger structure. Indeed,\nit is probably similar to the Western and Northeastern Clouds but at an even\nearlier point in its evolution. \\cite{Fesen92} note that the age of the\ninteraction is probably $4.1\\times 10^3$ years based on an assumed blast wave\nvelocity. Given the revised distance estimate of Blair et al. (1999), this age\nbecomes $2.3\\times 10^3$ years.\n\nIn \\Ha\\ (Figure~7a) we see a set of nonradiative filaments to the southeast of\nthe cloud. These filaments are visible very faintly in B5 (Figure~7c) as\nwell.Given the complete lack of X-ray emission (Figure~7d) to the east, these\nfilaments are the primary blast wave. The fact that these filaments are\nindented from the circular rim of the SNR implies the blast wave is diffracting\naround some object much larger than the visible emission and extended along our\nline of sight (\\cite{Graham95}).\n\nFesen et al. identify a filament segment seen to the west of the SE\ncloud--visible in both \\Ha\\ and our B5 image--as a reverse shock driven back\ninto the shocked medium. The X-ray emission, however, demonstrates that this\nis instead due the primary forward-moving blast wave. X-ray enhancement is\nseen to the west of the cloud, not the east as we would expect from a doubly\nshocked system. Furthermore, the optical filament is Balmer-dominated, which\nrequires a significant neutral fraction in the pre-shock gas, which would not\noccur at X-ray producing temperatures (\\cite{Graham95}). These points suggest\nthat the filament segment seen is a nonradiative piece of the main blast wave\nnot obviously related to the other emission in the area. The relative\nfaintness and lack of definition compared to other nonradiative filaments\nsuggests that it is not quite parallel to our line of sight.\n\nMeanwhile, the densest material in the shocked cloud tip has cooled enough to\nemit in ionic species like [\\ion{O}{3}] (Figure~7b) and \\ion{C}{4}. Gas\nstripping resulting from instabilities in the fluid flow along the edges of the\ncloud is seen as 'windblown streamers' on the north and south as well as\ndiffuse emission (because of a less favorable viewing angle) to the east. The\nB5 image shows great detail of the cloud shock and closely resemble the\n[\\ion{O}{3}] filaments, but with an added ``tail'' extending to the southeast.\nThe main body of the cloud shock as viewed in B5 is likely composed of\n\\ion{C}{4} and \\ion{O}{3}] emission while the ``tail'' may be an example of a\nslow shock in a neutral medium and have an enhanced two-photon flux\n(\\cite{Dopita82}). The shock velocity in the cloud is quoted by Fesen et al.\nas $<$60 \\vel\\ though this is based on the identification of the western\nsegment as a reverse shock. Given the bright [\\ion{O}{3}] and B5 emission in\nthe cloud shock, it seems more likely that the cloud shock is similar to other\nstructures to the north where shock velocities are thought to be more nearly\n140 \\vel.\n\nThere is a general increase in signal in the northern half of the SE FUV field\n(Figure~7c). It is unclear whether this is primarily due to the background\n``measles'' noted in \\S2 or if this represents diffuse, hot gas emitting\n\\ion{C}{4} as is seen in the halo around the central knot of XA. There is very\nfaint emission seen in both \\Ha\\ and [\\ion{O}{3}] in the area which could\nrepresent a region of more nearly face-on emitting gas.\n\n\\section{Concluding Remarks}\n\nThe UIT B5 band, although broader than ideal for SNR observations, provides a\nunique FUV spectral window. Under some conditions, the B5 bandpass provides\nimages of radiative filaments overrun by very high-speed shocks. Under other\nconditions, B5 observes nonradiative filaments at the extreme front edge of SNR\nblast waves. Combined with other image and spectral data, the B5 band can\nprovide unique insights into complex, difficult-to-model shock phenomena such\nas \\ion{C}{4} resonance scattering and shock completeness.\n\nIn nonradiative filaments, B5 flux comes from a mixture of \\ion{C}{4} as it\nionizes up and two-photon emission from preshock neutral hydrogen. In general,\nnonradiative filament morphology is very similar in B5 and \\Ha, implying that\ntwo-photon emission, originating in the same regions as \\Ha, is the primary\ncontributor to the B5 images. One unique capability of B5 imaging is its\nability to capture nonradiative shocks in ionized media. We see one example of\nsuch in Figure~6c where a nonradiative shock is faintly seen in \\ion{C}{4} and\n\\ion{He}{2}.\n\nRadiative filaments usually show good correlation between B5 and [\\ion{O}{3}]\nmorphology, suggesting that B5 flux arises in ions with similar excitation\nenergies such as \\ion{C}{4}. Existing models for simple, complete shocks\nindicate the same origin.\n\nHowever, existing FUV spectra complicate this picture, indicating that these\nregions should be dominated by two-photon flux which we would expect to follow\nmore closely the \\Ha\\ morphology. Observational selection restricts detailed\nspectral information to only the very brightest knots and filaments.\nPresumably, these bright regions also suffer the greatest resonance scattering\nin \\ion{C}{4}$\\lambda$1550, decreasing its observed flux; in fact, DBR found\nunexpectedly strong resonance scattering even away from the bright filaments\nand knots. Despite this, morphological similarities between B5 and\n[\\ion{O}{3}] in radiative filaments strongly suggest that, at least away from\nthe brightest filaments and cloud cores, B5 flux is dominated by \\ion{C}{4}.\n\n\\paragraph{Acknowledgements}\n\nThe authors wish to thank John Raymond for valuable discussions and the use of\nunpublished HUT data. We would also like to thank an anonymous referee for\nseveral valuable suggestions including using FUV images to trace nonradiative\nfilaments through ionized regions. Funding for the UIT project has been\nthrough the Spacelab Office at NASA headquarters under project number 440-551.\n\n\\begin{thebibliography}{Danforth00}\n\\bibitem[Arendt, Dwek, \\& Leisawitz 1992]{Arendt92}\nArendt, R.G., Dwek, E., \\& Leisawitz, D. 1992, \\apj, 400, 562\n\\bibitem[Bedogni \\& Woodward (1990)]{Bedogni90}\nBedogni, R. \\& Woodward, P. 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A., Binette, L. \\& Tuohy, I. R., 1984, \\apj, 282, 142\n\\bibitem[Fesen, Blair, \\& Kirshner 1982]{Fesen82}\nFesen, R. A., Blair, W. P., \\& Kirshner, R. P. 1982, \\apj, 262, 171\n\\bibitem[Fesen, Kwitter, \\& Downes (1992)]{Fesen92}\nFesen, R. A., Kwitter, K. B. \\& Downes, R. A. 1992, \\aj, 104, 719\n\\bibitem[Graham et al. 1995]{Graham95}\nGraham, J. R., Levenson, N. A., Hester, J. J., Raymond, J. C., \\& Petre, R.\n1995, \\apj, 444, 787\n\\bibitem[Green 1990]{Green90}\nGreen, D. A. 1990, \\aj, 100, 1927\n\\bibitem[HRH]{HRH}\nHartigan, P., Raymond, J. C., \\& Hartmann, L., 1987, \\apj, 316, 323 (HRH)\n\\bibitem[Hester \\& Cox 1986]{HesterCox86}\nHester, J. J., \\& Cox, D. P., 1986, \\apj, 300, 675\n\\bibitem[Hester 1987]{Hester87}\nHester, J. J. 1987, \\apj, 314, 187\n\\bibitem[Hester, Raymond, \\& Danielson 1986]{Hester86}\nHester, J. J., Raymond, J. C., \\& Danielson, G. E. 1986, \\apj, 303, L17\n\\bibitem[Hester, Raymond, \\& Blair 1994]{Hester94}\nHester, J. J., Raymond, J. C., \\& Blair, W. P. 1994, \\apj, 420, 721\n\\bibitem[Innes (1992)]{Innes92}\nInnes, D. E. 1992, \\aap, 256, 660\n\\bibitem[Leahy et al. 1997]{Leahy97}\nLeahy, D. A., Roger, R. S., \\& Ballantyne, D. 1997, \\aj, 114, 2081\n\\bibitem[Levenson et al. 1996]{Levenson96}\nLevenson, N. A., Graham, J. R., Hester, J. J., \\& Petre, R. 1996, \\apj, 468,\n323\n\\bibitem[Levenson et al. 1997]{Levenson97}\nLevenson, N. A., et al. 1997, \\apj, 484, 304\n\\bibitem[Levenson et al. 1998]{Levenson98}\nLevenson, N. A., Graham, J. R., Keller, L. D., \\& Richter, M. J. 1998, \\apjs,\n118, 541\n\\bibitem[Long et al. (1992)]{Long92}\nLong, K. S., et al. 1992, \\apj, 400, 214\n\\bibitem[Miller (1974)]{Miller74}\nMiller, J. S. 1974, \\apj, 189, 239\n\\bibitem[Nussbaumer \\& Schmutz 1984]{Nussbaumer84}\nNussbaumer, H., \\& Schmutz, W. 1984, \\aap 138, 495\n\\bibitem[Osterbrock 1989]{Osterbrock89}\nOsterbrock, D. S., 1989, ``\\it{Astrophysics of Gaseous Nebulae and Active\nGalactic Nuclei}\\rm'', Mill Valley, CA, University Science Books\n\\bibitem[Raymond (1979)]{Raymond79}\nRaymond, J. C. 1979, \\apjs, 39, 1\n\\bibitem[Raymond et al. (1980a)]{Raymond80a}\nRaymond, J. C., Davis, M., Gull, T. R., \\& Parker, R. A. R. 1980a, \\apj, 238,\nL21\n\\bibitem[Raymond et al. (1980b)]{Raymond80b}\nRaymond, J. C., Black, J. H., Dupree, A. K., Hartmann, L., \\& Wolff, R. S.\n1980b, \\apj, 238, 881\n\\bibitem[Raymond et al. 1981]{Raymond81}\nRaymond, J. C., Black, J. H., Dupree, A. K., Hartmann, L., \\& Wolff, R. S.\n1981, \\apj, 246, 100\n\\bibitem[Raymond et al. 1983]{Raymond83}\nRaymond, J. C., Blair, W. P., Fesen, R. A. \\& Gull, T. R. 1983, \\apj, 324, 869\n\\bibitem[Raymond et al. 1988]{Raymond88}\nRaymond, J. C., et al. 1988, \\apj, 324, 869\n\\bibitem[Scoville et al. 1977]{Scoville77}\nScoville, N. Z., Irvine, W. M., Wannier, P. G., \\& Predmore, C. R. 1977, \\apj,\n216, 320\n\\bibitem[Smith et al. 1996]{Smith96}\nSmith, E. P., et al. 1996, \\apjs, 104, 287\n\\bibitem[Stecher et al. (1992)]{Stecher92}\nStecher, T. P., et al. 1992, \\apj, 395, L1\n\\bibitem[Stecher et al. 1997]{Stecher97}\nStecher, T. P., et al. 1997, \\pasp, 109, 584\n\\bibitem[Stone \\& Norman (1992)]{Stone92}\nStone, J. M., \\& Norman, M. L., 1992, \\apjl, 390, L17\n\\bibitem[Szentgyorgyi et al. 2000]{Szentgyorgyi99}\nSzentgyorgyi, A. H., Raymond, J. C., Hester, J. J., \\& Curiel, S. 2000, \\apj,\nin press\n\\bibitem[Treffers 1981]{Treffers81}\nTreffers, R. R. 1981, \\apj, 250, 213\n\\end{thebibliography}\n\n\\clearpage\n\n\\begin{figure}\n\\centerline{\\psfig{figure=fig1.ps,width=6in}}\n\\figcaption[]{\\protect\\small{Two typical UV spectra of SNR filaments. The top, a\nradiative filament (Blair et al. 1991), shows lines of many different ions.\nThe bottom, a nonradiative filament (Long et al. 1992), shows lines of only the\nhighest ionization species. The dashed curve superimposed on the two spectra\nrepresents the throughput of UIT's B5 filter as a function of wavelength.}}\n\\end{figure}\n\n\\begin{figure}\n\\centerline{\\psfig{figure=fig2.ps,width=5in}}\n\\figcaption[]{\\protect\\small{An \\Ha\\ mosaic of the entire Cygnus Loop (courtesy\nLevenson et al. 1998) with the 40\\arcmin\\ UIT fields superimposed and\nlabeled.}}\n\\end{figure}\n\n\\begin{figure}\n\\figcaption[]{\\protect\\small{The Western Cloud as viewed in a) \\Ha, b)\n[\\protect\\ion{O}{3}]$\\lambda$5007, c) B5, and d) the ROSAT High Resolution Imager\n(HRI). All bands show a bright north-south radiative complex similar to that\nseen in the Northeast Cloud (Figure~5) but with apparently simpler geometry.\nTwo bright, parallel filaments suggest two points of tangency to our line of\nsight. In \\Ha\\ (a) and B5 (c) we also see a nonradiative filament diverging to\nthe northwest of the bright radiative region. X-rays are seen behind this\nfilament in (d). A reverse shock generates higher temperatures and brighter\nX-ray emission at the radiative region. With the exception of the nonradiative\nfilament, the B5 and the [\\protect\\ion{O}{3}] (b) show a high degree of correlation,\nsuggesting origin of the B5 flux in high-excitation ionic species. Each field\nis 40\\arcmin\\ across. For Figures 3-7, both optical fields have been\nmedian-filtered with a 3-pixel (15 \\arcsec) box. The HRI field has been\nsmoothed with a 3-pixel gaussian. All fields are aligned and oriented with\nnorth at the top and east to the left. All image intensities are displayed\nlogarithmically. (Please see attached file uitfig3.jpg or \\url\\ for\nfull-resolution image.)}}\n\n\\figcaption[]{\\protect\\small{The Northeast Nonradiative Region, containing the classic\nnonradiative filaments, viewed as in Figure~3. The SN blast wave propagates\nthrough the atomic shell at v$\\sim$400 \\vel. Thin filamentary emission arises\nfrom the preshock neutral fraction as it heats up, and is seen in \\Ha\\ (a).\nLittle or no emission is seen from this filament in [\\protect\\ion{O}{3}] (b). The\nshock is visible in B5 (c) through both two-photon processes (closely linked to\n\\protect\\Ha\\ emission) and to a lesser extent through high-ionization species--in this\ncase \\protect\\ion{C}{4}. The ROSAT HRI image (d) shows the $\\sim10^{6}$K\nX-ray-emitting post-shock gas in a band behind the shock front. (Please see\nattached file uitfig4.jpg or \\url\\ for full-resolution image.)}}\n\n\\figcaption[]{\\protect\\small{The Northeast Cloud as viewed in Figure~3. \\protect\\Ha\\ (a) shows\nsmooth nonradiative filaments to the east of a more complex mass of radiative\nfilaments. [\\protect\\ion{O}{3}] (b) shows a radiative filament structure complicated by\nline-of-sight coincidence of several emitting regions. The short B5 exposure\n(c) shows radiative structures well but is not deep enough to show the\nnonradiative filaments. The ROSAT image (d) shows that flux from hot gas\nbehind the blast wave is considerably enhanced by the strongly decelerating\nshocks in the denser radiative region. (Please see attached file uitfig5.jpg\nor \\url\\ for full-resolution image.)}}\n\n\\figcaption[]{\\protect\\small{The XA region as viewed in Figure~3. This region displays\na complex region of cloud-shock interactions, including dense, bright filaments\nwhose \\protect\\ion{C}{4} emission is apparently strongly affected by resonance\nscattering and a range of shock completeness (see Figures 9c and 10). (Please\nsee attached file uitfig6.jpg or \\url\\ for full-resolution image.)}}\n\n\\figcaption[]{\\protect\\small{The South East Cloud viewed as in Figure~3. This small\npatch of emission is likely the tip of a much larger cloud early in the stages\nof shock interaction. B5 morphology (c) closely matches both \\protect\\Ha\\ (a) and\n[\\protect\\ion{O}{3}] (b). The primary differences are faint features common only to B5\nand \\Ha: the ``tail'' extending south of the cloud, and faint diffuse\nmaterial--nonradiative blast wave filaments--in the southern half of the field.\n X-rays (d) show a region of emission much larger than the optical/UV cloud\nwith a slight flux depression or ``hole'' in the center. (Please see attached\nfile uitfig7.jpg or \\url\\ for full-resolution image.)}}\n\n\\figcaption[]{\\protect\\small{Three-color images of the Western Cloud (a), the Northeast\nCloud (b), the XA Region (c) and the Southeast Cloud (d). \\protect\\Ha\\ is in red, B5\nemission in green and X-rays in blue. All intensities are displayed\nlogarithmically and colors have been adjusted to best show spatial\nrelationships. (Please see attached file uitfig8.jpg or \\url\\ for\nfull-resolution image.)}}\n\\end{figure}\n\n\\begin{figure}\n\\centerline{\\psfig{figure=fig9.ps,width=6.5in}}\n\\figcaption[]{\\protect\\small{The effects of completeness and resonance scattering. a)\nUIT B5 image of the XA region. b) A ``completeness map'' generated by taking\nthe ratio of B5 to \\protect\\Ha. Light areas represent more complete cooling while dark\nareas are less complete. Complete shocks emit predominantly two-photon\nemission in the B5 band while incomplete regions tend toward higher \\protect\\ion{C}{4}\ncontributions. c) ``Saturation map'' generated from B5 and [\\protect\\ion{O}{3}]. Dark\nregions show areas of higher resonance scattering of \\protect\\ion{C}{4}.}}\n\\end{figure}\n\n\\begin{figure}\n\\centerline{\\psfig{figure=fig10.eps,width=3.5in}}\n\\figcaption[]{\\protect\\small{B5 flux as a function of velocity for the shock models of\nHartigan, Raymond, \\& Hartmann (1987). The various dashed lines show the flux\nfrom selected ionic species multiplied by the B5 filter throughput, as a\nfunction of radiative shock velocity. The solid line is the total B5 flux\ncalculated as a sum of 0.80$\\times$\\protect\\ion{C}{4}$\\lambda$1550,\n0.55$\\times$\\protect\\ion{He}{2}$\\lambda$1640, 0.54$\\times$\\protect\\ion{O}{3}]$\\lambda$1665, and\n0.15$\\times$two-photon emission. These models hold only for complete,\nsingle-velocity shocks and do not take account of resonance scattering or other\ncomplications. In these models \\protect\\ion{C}{4} emission dominates the B5 bandpass\nat all but the lowest velocities.}}\\label{models}\n\\end{figure}\n\n\\end{document}\n\n\n\n\n\n\n\n\n\n\n"
}
] |
[
{
"name": "astro-ph0002005.extracted_bib",
"string": "\\begin{thebibliography}{Danforth00}\n\\bibitem[Arendt, Dwek, \\& Leisawitz 1992]{Arendt92}\nArendt, R.G., Dwek, E., \\& Leisawitz, D. 1992, \\apj, 400, 562\n\\bibitem[Bedogni \\& Woodward (1990)]{Bedogni90}\nBedogni, R. \\& Woodward, P. R. 1990, \\aap, 231, 481\n\\bibitem[Benvenuti, Dopita, \\& D'Odorico (1980)]{Benvenuti80}\nBenvenuti, P., Dopita, M., \\& D'Odorico, S. 1980, \\apj, 238, 601\n\\bibitem[Blair et al. 1991]{Blair91}\nBlair, W. P., et al. 1991, \\apjl, 379, L33\n\\bibitem[Blair et al. 1999]{Blair99}\nBlair, W. P., Sankrit, R., Raymond, J. C. \\& Long, K. S., 1999, AJ, 118, 942\n\\bibitem[Chevalier, Kirshner, \\& Raymond 1980]{CKR80}\nChevalier, R. A., Kirshner, R. P., \\& Raymond, J. C. 1980, \\apj, 235, 186\n\\bibitem[Chevalier \\& Raymond 1978]{Chevalier78}\nChevalier, R. A., \\& Raymond, J. C. 1978, \\apjl, 225, L27\n\\bibitem[Cornett et al. 1992]{Cornett92}\nCornett, R. H., et al. 1992, \\apj, 395, L9\n\\bibitem[DBR]{DBR}\nDanforth, C. W., Blair, W. P., \\& Raymond, J. C. 2000, in prep. (DBR)\n\\bibitem[Dopita, Binette, \\& Schwartz, 1982]{Dopita82}\nDopita, M. A., Binette, L., \\& Schwartz, R. D. 1982, \\apj, 261, 183\n\\bibitem[Dopita, Binette, \\& Tuohy (1984)]{Dopita84}\nDopita, M. A., Binette, L. \\& Tuohy, I. R., 1984, \\apj, 282, 142\n\\bibitem[Fesen, Blair, \\& Kirshner 1982]{Fesen82}\nFesen, R. A., Blair, W. P., \\& Kirshner, R. P. 1982, \\apj, 262, 171\n\\bibitem[Fesen, Kwitter, \\& Downes (1992)]{Fesen92}\nFesen, R. A., Kwitter, K. B. \\& Downes, R. A. 1992, \\aj, 104, 719\n\\bibitem[Graham et al. 1995]{Graham95}\nGraham, J. R., Levenson, N. A., Hester, J. J., Raymond, J. C., \\& Petre, R.\n1995, \\apj, 444, 787\n\\bibitem[Green 1990]{Green90}\nGreen, D. A. 1990, \\aj, 100, 1927\n\\bibitem[HRH]{HRH}\nHartigan, P., Raymond, J. C., \\& Hartmann, L., 1987, \\apj, 316, 323 (HRH)\n\\bibitem[Hester \\& Cox 1986]{HesterCox86}\nHester, J. J., \\& Cox, D. P., 1986, \\apj, 300, 675\n\\bibitem[Hester 1987]{Hester87}\nHester, J. J. 1987, \\apj, 314, 187\n\\bibitem[Hester, Raymond, \\& Danielson 1986]{Hester86}\nHester, J. J., Raymond, J. C., \\& Danielson, G. E. 1986, \\apj, 303, L17\n\\bibitem[Hester, Raymond, \\& Blair 1994]{Hester94}\nHester, J. J., Raymond, J. C., \\& Blair, W. P. 1994, \\apj, 420, 721\n\\bibitem[Innes (1992)]{Innes92}\nInnes, D. E. 1992, \\aap, 256, 660\n\\bibitem[Leahy et al. 1997]{Leahy97}\nLeahy, D. A., Roger, R. S., \\& Ballantyne, D. 1997, \\aj, 114, 2081\n\\bibitem[Levenson et al. 1996]{Levenson96}\nLevenson, N. A., Graham, J. R., Hester, J. J., \\& Petre, R. 1996, \\apj, 468,\n323\n\\bibitem[Levenson et al. 1997]{Levenson97}\nLevenson, N. A., et al. 1997, \\apj, 484, 304\n\\bibitem[Levenson et al. 1998]{Levenson98}\nLevenson, N. A., Graham, J. R., Keller, L. D., \\& Richter, M. J. 1998, \\apjs,\n118, 541\n\\bibitem[Long et al. (1992)]{Long92}\nLong, K. S., et al. 1992, \\apj, 400, 214\n\\bibitem[Miller (1974)]{Miller74}\nMiller, J. S. 1974, \\apj, 189, 239\n\\bibitem[Nussbaumer \\& Schmutz 1984]{Nussbaumer84}\nNussbaumer, H., \\& Schmutz, W. 1984, \\aap 138, 495\n\\bibitem[Osterbrock 1989]{Osterbrock89}\nOsterbrock, D. S., 1989, ``\\it{Astrophysics of Gaseous Nebulae and Active\nGalactic Nuclei}\\rm'', Mill Valley, CA, University Science Books\n\\bibitem[Raymond (1979)]{Raymond79}\nRaymond, J. C. 1979, \\apjs, 39, 1\n\\bibitem[Raymond et al. (1980a)]{Raymond80a}\nRaymond, J. C., Davis, M., Gull, T. R., \\& Parker, R. A. R. 1980a, \\apj, 238,\nL21\n\\bibitem[Raymond et al. (1980b)]{Raymond80b}\nRaymond, J. C., Black, J. H., Dupree, A. K., Hartmann, L., \\& Wolff, R. S.\n1980b, \\apj, 238, 881\n\\bibitem[Raymond et al. 1981]{Raymond81}\nRaymond, J. C., Black, J. H., Dupree, A. K., Hartmann, L., \\& Wolff, R. S.\n1981, \\apj, 246, 100\n\\bibitem[Raymond et al. 1983]{Raymond83}\nRaymond, J. C., Blair, W. P., Fesen, R. A. \\& Gull, T. R. 1983, \\apj, 324, 869\n\\bibitem[Raymond et al. 1988]{Raymond88}\nRaymond, J. C., et al. 1988, \\apj, 324, 869\n\\bibitem[Scoville et al. 1977]{Scoville77}\nScoville, N. Z., Irvine, W. M., Wannier, P. G., \\& Predmore, C. R. 1977, \\apj,\n216, 320\n\\bibitem[Smith et al. 1996]{Smith96}\nSmith, E. P., et al. 1996, \\apjs, 104, 287\n\\bibitem[Stecher et al. (1992)]{Stecher92}\nStecher, T. P., et al. 1992, \\apj, 395, L1\n\\bibitem[Stecher et al. 1997]{Stecher97}\nStecher, T. P., et al. 1997, \\pasp, 109, 584\n\\bibitem[Stone \\& Norman (1992)]{Stone92}\nStone, J. M., \\& Norman, M. L., 1992, \\apjl, 390, L17\n\\bibitem[Szentgyorgyi et al. 2000]{Szentgyorgyi99}\nSzentgyorgyi, A. H., Raymond, J. C., Hester, J. J., \\& Curiel, S. 2000, \\apj,\nin press\n\\bibitem[Treffers 1981]{Treffers81}\nTreffers, R. R. 1981, \\apj, 250, 213\n\\end{thebibliography}"
}
] |
astro-ph0002006
|
ULTRA HIGH ENERGY COSMIC RAYS: the theoretical challenge
|
[
{
"author": "A. V. Olinto\\thanksref{corr}"
}
] | "The origin of the highest-energy cosmic rays remains a mystery. The lack of a high energy cutoff in(...TRUNCATED) | [{"name":"physrepf.tex","string":"\\def\\la{\\hbox{{\\lower -2.5pt\\hbox{$<$}}\\hskip -8pt\\raise\n-(...TRUNCATED) | [{"name":"astro-ph0002006.extracted_bib","string":"\\begin{thebibliography}{9}\n\n\\bibitem{watson99(...TRUNCATED) |
astro-ph0002007
|
[] | [{"name":"astro-ph0002007.tex","string":"\\documentstyle[12pt]{article}\n\\textheight 9in\n\\textwid(...TRUNCATED) | [{"name":"astro-ph0002007.extracted_bib","string":"\\begin{thebibliography}{99}\n\\bibitem{alcock} A(...TRUNCATED) |
||
astro-ph0002008
| "ISO-SWS spectroscopy of NGC 1068 \\footnote{Based on observations with ISO, an ESA project with ins(...TRUNCATED) | [{"author":"D.~Lutz\\footnote{Max-Planck-Institut f\\\"ur extraterrestrische Physik, Postfach 1603, (...TRUNCATED) | "We present ISO-SWS spectroscopy of NGC 1068 for the complete wavelength range 2.4 to 45$\\mu$m at r(...TRUNCATED) | [{"name":"obspaper.tex","string":"%\\documentstyle[12pt,aasms4]{article}\n\\documentstyle[11pt,aaspp(...TRUNCATED) | [{"name":"astro-ph0002008.extracted_bib","string":"\\begin{thebibliography}{}\n\\bibitem[Alexander e(...TRUNCATED) |
astro-ph0002009
|
The giant radio galaxy 8C\,0821+695 and its environment
| [{"author":"L. Lara\\inst{1}"},{"author":"K.-H. Mack\\inst{2,3}"},{"author":"M. Lacy\\inst{4}"},{"au(...TRUNCATED) | "We present new VLA and Effelsberg observations of the radio galaxy 8C\\,0821+695. We have obtained (...TRUNCATED) | [{"name":"lara.tex","string":"%----------------------------------------------%\n% REFERENCES \n%----(...TRUNCATED) | [{"name":"astro-ph0002009.extracted_bib","string":"\\begin{thebibliography}{}\n\n\\bibitem[1977]{baa(...TRUNCATED) |
astro-ph0002010
|
RADIATION SPECTRA FROM \\ADVECTION-DOMINATED ACCRETION FLOWS \\IN A GLOBAL MAGNETIC FIELD
|
[
{
"author": "MOTOKI KINO$^{1}$"
},
{
"author": "OSAMU KABURAKI and NAOHIRO YAMAZAKI"
}
] | "We calculate the radiation spectra from advection-dominated accretion flows (ADAFs), taking into ac(...TRUNCATED) | [{"name":"ms.tex","string":" \n%--------------------------------------------------------------------(...TRUNCATED) |
[] |
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