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