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Mapping the intergalactic medium in emission
at low and high redshiftSerena Bertone (UC Santa Cruz)Joop Schaye (Leiden Observatory)
&
the OWLS Team
Queen’s University, KingstonApril 3rd 2009
Outline
Introduction: the intergalactic medium: warm, warm-hot, hot… what is it?
how do we detect it?
New simulations: OWLS
Results: X-ray & UV emission at z<1
rest-frame UV emission at z>2
what gas does emission trace?
dependence on physics
unbiased info on matter power spectrum on largest scales
reservoir of fuel for star and galaxy formation
IGM metallicity constrains the cosmic SF history
interplay IGM-feedback puts constraints on galaxy formation models
Intergalactic medium: Why do we care?
Neutrinos:0.3%
Metals: 0.03%
Dark energy: 70%
Free H & He: 4%
Stars: 0.5%
Dark matter: 25%
> 90% of the baryonic mass is in diffuse gas(Persic & Salucci 1992, Fukugita et al 1998, Cen & Ostriker 1999)
100 Mpc
z=6 z=2 z=0
Springel 2003
IGM evolution
cosmic time
Gas temperature distribution
Cen & Ostriker 1999 etc
T<105 K 105 K <T<107 K T>107 K
Warm gas:
diffused
photo-ionised
traced by Lyα forest
Warm-hot gas: WHIM mildly overdense
collisionally ionised
shock heated by gravitational shocks
Hot gas: dense metal enriched haloes of galaxies and
clusters
z=0
IGM mass and metals
The bulk of the metal mass does not trace the bulk of the IGM mass
most mass
mostmetalmass
Bertone et al. 2009a
z=0.25
Mass and metal fractions evolution
Average IGM temperature increases with time
Most metals locked in stars at z=0
Average metal temperature increases with time
Metal mass - Wiersma et al 2009b
WHIMICM
diffuse IGM
halo gas
IGM mass - Dave’ et al. 2001
T<106 KRest-frame UV linesLyα, OVI, CIV…
T>106 KSoft X-rays metal lines
OVII, OVIII, FeXVII
Absorption
z<1Nicastro et al 2005
Rasmussen et al 2007Buote et al 2009
z<1: UVTripp et al 2007Lehner et al 2007Danforth & Shull 2005
z>1.5: opticalKim et al 2001Simcoe et al 2006
z<1Fang et al 2005Bertone et al 2009a
z<1: UVFurlanetto et al 2004 Bertone et al 2009a
z>1.5: opticalWeidinger et al 2004Bertone et al 2009b
Emission
how can we detect the igm?
✔✔
✗✔✗
✗
✔?
Emission vs. absorption
Emission Absorption
PRO• 3-D data on spatial distribution, velocity field, metallicity…
PRO• Easier to detect• Underdense regions investigated with current optical/UV instruments
CONS• Hard to detect – low fluxes• Requires new instruments with high sensitivity and large FoV
CONS• 1-D info along a LOS• Requires bright background sources
Owls
OverWhelmingly Large Simulations
The OWLS Team:
Joop Schaye (PI), Claudio Dalla Vecchia, Rob Wiersma,Craig Booth, Marcel Haas (Leiden)Volker Springel (MPA Garching)Luca Tornatore (Trieste)Tom Theuns (Durham)+ the Virgo Consortium
Many thanks to the LOFAR and SARA supercomputing facilities
Owls
cosmological hydrodynamical simulations: Gadget 3
many runs with varying physical prescriptions/numerics
run on LOFAR IBM Bluegene/L
WMAP 3 cosmology
largest runs: 2x5123 particles
two main sets: L=25 Mpc/h and L=100 Mpc/h boxes
evolution from z>100 to z=2 or z=0
New physics in Owls
New star formation (Schaye & Dalla Vecchia 2008): Kennicutt-Schmidt SF law implemented without free parameters
New wind model (Dalla Vecchia & Schaye 2008): winds local to the SF event hydrodynamically coupled
Added chemodynamics (Wiersma et al. 2009): 11 elements followed explicitly (H, He, C, N, O, Ne, Si, Mg, S, Ca, Fe) Chabrier IMF SN Ia & AGB feedback
New cooling module (Wiersma, Schaye & Smith 2009): cooling rates calculated element-by-element from 2D CLOUDY tables photo-ionisation by evolving UV background included
Physics variations in OWLS
Cosmology: WMAP1 vs WMAP3 vs WMAP5Reionisation & Helium reionisationGas cooling: primordial abundances vs metal dependentStar formation: top heavy IMF in bursts isothermal & adiabatic EoS Schmidt law normalisation Metallicity-dependent SF thresholds
Feedback: no feedback feedback intensity: mass loading, initial velocity… feedback implementation (Springel & Hernquist; Oppenheimer & Dave’) AGN feedback
Chemodynamics: Chabrier vs Salpeter IMF SN Ia enrichment AGB mass transfer
Gas cooling rates
Wiersma, Schaye & Smith 2008
collisionalionisation eq.
photoionisation eq.density dependent
Photo-ionisation by UV BK + collisional ionisation equilibrium
Cooling rates calculated element by element for 11 species: takes into account changes in the relative abundances
Gas emissivity
11 elements – comparable to cooling rates
Collisional ionisation + photo-ionisation by UV BK important at low density
UV lines stronger than X-ray ones
UV lines X-ray lines
Den
sity
Bertone et al 2009a
z=0.25
Emission at low redshift:
Soft X-rays& UV lines
Emission at low redshift: X-rays
100 Mpc/h boxes20 Mpc/h thick slices
15” angular resolution12 emission lines
Bertone et al 2009a
X-ray lines
O VIII strongest line
lines from lower ionisation states and whose emissivity peaks at lower temperatures trace moderately dense IGM: C V, C VI, N VII, O VII, O VIII and Ne IX
lines from higher ionisation states trace denser, hotter gas: C VI, O VIII, Ne X, Mg XII, Si XIII, S XV and Fe XVII
Fe XVII emission has different spatial distribution than other elements: later enrichment by SN Ia
Bertone et al 2009a
What gas produces x-ray emission?
X-ray emission
traceswarm-hot, moderately dense gas,
not the bulk of the IGM mass and metals.
emission-weighted particle distributions
Bertone et al 2009a
Emission at low redshift: UV
Bertone et al 2009a
UV emission is a good tracer of galaxies and of mildly dense IGM,but not of IGM in very dense environments
UV lines
C IV strongest line: traces gas in proximity of galaxies
different spatial distribution: O VI and Ne VIII trace more diffuse gas than C IV
no emission from the hottest gas in groups
What gas produces UV emission?emission-weighted particle distributions
Bertone et al 2009a
O VI and Ne VIII trace WHIM gas
C IV traces more metal-rich, cooler gas
Summary:Dependencies on gas properties
Correlation of emission with density and metallicity: highest emission from densest and most metal enriched particles
Median temperature of highest emission corresponds to peak temperature of emissivity curve – as seen in T-nHI diagrams
Median density, temperature and metallicity of gas particles vs. particle emission
Bertone et al 2009a
Impact of physics
What happens when changing the physical model?
Impact of physics: x-rays
no feedback: no metal transport localised emission
primordial cooling rates: longer cooling times stronger emission at high density (≈100 times)
momentum driven winds: metals more spread in IGM weaker emission (≈100 times)
AGN feedback: weaker emission in dense regions
Bertone et al 2009a
no feedback: emission localised in galaxies
primordial cooling rates: stronger emission at high density
AGN feedback: weaker emission
changes in wind parameters: small effect
Impact of physics: UV
Bertone et al 2009a
Emission at high redshift:rest-frame UV lines
Bertone et al 2009b
emission at 2<z<5
25 Mpc/h simulations
2” angular resolution
IGM emission at z>1.5
At z>1.5 rest-frame UV lines are redshifted in to the optical band
A number of upcoming optical instrument might detect IGM emission lines at 1.5<z<5:
Cosmic Web Imager on Palomar (CWI, Rahman et al 2006) this year!
Keck Cosmic Web Imager (KCWI)
Antarctic Cosmic Web Imager (ACWI, Moore et al 2008)
MUSE on VLT (Bacon et al 2009)
IFUs with large fields of view & high spatial resolution
Great chance to observe the 3-D structure of the IGM for the first time!
Hig
her i
onis
ation
stat
es
Low
er io
nisa
tion
stat
es
Emission PDFs
Lower ionisation states: single lines shorter wavelengths C III up to 10x stronger than C IV
Higher ionisation states: doublets – easy to identify weaker than lower ion. states
What gas produces UV emission?emission-weighted particle distributions
C and Si lines trace colder, more metal enriched gas than O, N and Ne lines Most emission comes from dense gas more metal enriched More highly ionised lines trace warmer gas (e.g. C IV vs C III)
What can we observe with CWI?
CWI: flux limit: 100
photon/s/cm2/sr angular resolution: 2”
Everything in red and white might be detectable
central regions of groups and galactic haloes
z=2
Lines detectable by CWI
Line Wavelength(Å)
Minimum redshift Detectable to z
C III 977.0 2.7 ≤ 4
C IV 1548 1.4 ≤ 3
N IV 765 3.7 ×
N V 1230.9 1.9 ≤ 2
O IV 549.3 5.6 ×
O V 630.0 4.7 ×
O VI 1032 2.5 ≤ 3
Ne VIII 770.0 3.7 ×
Si III 1207 2.0 ≤ 3
Si IV 1394 1.6 ≤ 3
summaryDense cool gas in the haloes of galaxies is traced by:• UV lines: C IV, Si IV
Low density WHIM gas in filaments is traced by:• UV lines: O VI, Ne VIII
• soft X-ray lines from hydrogen-like atoms: C V, O VII, Ne IX
• soft X-ray lines from elements with low atomic numbers: OVIII
Dense hot gas in clusters and groups is traced by:• soft X-ray lines from fully ionised atoms: C VI, OVIII
• soft X-ray lines from elements with high atomic numbers: Mg XII, Fe XVII
Detection of WHIM emission by future telescopes:• challenging in low density regions
• very likely in groups and cluster outskirts
• CWI very likely to detect metal line emission at high z for the first time
Convergence tests
Angular resolution:needed for detecting the spatial distribution of the emitting gas
Number of particles:need high resolutionfor physicalprocesses(star formation,feedback, coolingetc) to converge
Emission per unit volume
Bertone et al 2009b
Most energy emitted by lower ionisation lines: C III, O IV, O V
O VI strongest line at higher ionisation
Lot of energy emitted by Si lines, despite low emissivity
Ne and N emission per unit volume very weak
strongest X-ray emission from dense regionsweak emission from the low density, metal poor gas
Emission vs density
Density cut maps