X-ray Studies of Interstellar and Intergalactic Dustlia/Corrales_grad_colloquium.pdfX-ray scattering is a diagnostic tool for ISM structure ... D. Russeil: Star-forming complexes and

X-ray Studies of Interstellar and

Intergalactic Dust

Lia CorralesColumbia UniversityNASA Earth and Space Science Fellow

Advised by Frits Paerels

X-ray Studies of ISM and IGM Dust

Lia Corrales - CU Colloquium - May 7, 2014

X-ray scattering tools for studying the ISM

Cygnus X-3:

Grain sizes and spatial distribution

Dust-to-gas mass ratio

Dust in the intergalactic medium:

Future prospects



Dust extinction (MRN)

X-ray Optical J,H,K,L Herschel

Spitzer



Optical (DSS) Infrared (SFD) X-ray (RASS)

Credit: World Wide Telescope











Cygnus X-3 (Chandra)

dust



X-ray scattering is a diagnostic tool for ISM grain sizes

SGR J1550-5418(NASA/Swift/Halpern)

Strongly forward (small angle) scattering

10

a(µm) E(keV)

Strongly sensitiveto grain size sca / a4E2

aX-ray light


Lia Corrales - AAS 223, Jan 2014

Strongly forward (small angle) scattering

10

a(µm) E(keV)

Strongly sensitiveto grain size sca / a4E2

aX-ray light

X-ray scattering is a diagnostic tool for judging distance




aX-ray light

↵

sca

↵

OS

X-ray scattering is a diagnostic tool for ISM structure




↵

sca

↵

OS

Screen case





↵

sca

↵

OS

Screen case





↵

sca

↵

OS

Screen case





sca

↵

OS

Uniform case


↵

↵




Cygnus X-3:




Future prospects



Assume dN

da/ ap

Fit foramax

p

Md

Cut-off grain size:

Power law exponent:

Dust mass column:

Cyg X-3 Scattering Halo [1-6 keV]

Corrales & Paerels (2014)



Cyg X-3 Scattering Halo [1-6 keV]




Bayesian analysis (emcee) finds a population of likely fits


sca(1 keV) = 2.4

amax

= 0.14 µmp = 3.5

[2 = 1.3]



D. Russeil: Star-forming complexes and the spiral structure of our Galaxy 143

Fig. 5. The adopted four-arm model. The symbol size is proportional to the excitation parameter. The Sun position is given by the large starsymbol. All the complexes are plotted. 1: Sagittarius-Carina arm, 2: Scutum-Crux arm, 10: Norma-Cygnus arm and 20: Perseus arm. We havealso sketched the local arm feature (long dashed line), the bar orientation and length (dashed-dot-dot line) from Englmaier & Gerhard (1999),the expected departure from a logarithmic spiral arm observed for the Sagittarius-Carina arm (short dashed line) and finally a feature certainlylinked to the three-kpc arm (solid line)

As we already noticed we have fitted regular logarithmicarms, but it is commonly seen on external galaxies that arms donot have a regular design. This fact is illustrated here because aprecise study of the nearest parts of the Sagittarius-Carina armplaces it at 2 kpc from the Sun while the fitted spiral passes at1 kpc: the axis of the arm must be moved slightly inward inregions just inside the Sun (Fig. 5, dashed line).

The four arm model is consistent with the recent work ofAmaral & Lepine (1997) from a dynamical approach. Theyshow that the spiral structure of our Galaxy can be representedby the superposition of two- and four-arm components witha pitch angle of about 14; such a solution looks like a purefour-arm structure when the components are in phase. Thisis another argument in favor of the four-arm model (Fig. 5).Comparing the adopted model to the model of Georgelin &Georgelin (1976) one can see that the general design is pre-served but now the arms are more clearly defined and moreextended. The Norma segment appears now to be connected to

the Cygnus segment, forming a single extended arm symmetri-cal to the Sagittarius-Carina arm. The Scutum-Crux arm, whichis the shortest, is symmetric with respect to the Perseus arm.The Carina arm appears to be a major arm; sometimes such amajor arm is observed in grand design external galaxies. Thissuggests a grand design structure coexisting with smaller-scalestructures, rather than a purely floculent one.

One of these smaller-scale structures is the local arm.Figure 5 (long dashed line) shows a hint of the presence ofthe local arm: we note a collection of aligned complexes paral-lel to the Sagittarius arm and including the Sun. Unfortunately,from our data it is impossible to define its characteristics. Jacqet al. (1988) have already showed the existence of a collectionof molecular clouds aligned with a pitch angle of 22 and link-ing up the local material to the Sagittarius arm.

Three arms out of four start close to the end of the Galacticbar major axis. In barred galaxies, arms are generally con-nected to the bar (Sellwood & Sparke 1988). Moreover, for our

Perseus

Sagittarius

Norma-Cygnus

Cyg X-3

l = 79.8b = +00.7

[Russeil 2003]

Milky Way spiral structure as probed by star forming complexes


Lia Corrales - CU Colloquium - May 7, 2014Corrales & Paerels (2014)


12



12

Screen 2:20% of masswithin 1 kpcof Cyg X-3

Screen 1:80% of mass

consistent withPerseus arm



[2 = 2.8]

amax

= 0.21 µmp = 3.6

sca(1 keV) = 1.9




Cygnus X-3:




Future prospects



X-ray spectral fitting yields a dust-to-gas mass ratio


sca(1 keV) = 2.2



X-ray spectral fitting yields a dust-to-gas mass ratio


dust-to-gasmass ratio

MilkyWay 1/3 4/3dust-to-gas

mass ratioMilkyWay 1/2 4/3



X-ray scattering and absorption gauges:




grain size distribution





ISM structure





ISM structure

distance





ISM structure

distance

dust-to-gas ratio





ISM structure

distance

dust-to-gas ratio

dust-to-metal ratio





ISM structure

distance

dust-to-gas ratio

dust-to-metal ratio

dust composition




Cygnus X-3:




Future prospects



1. Could intergalactic dust interfere with dark energy surveys?




of information on the shapes of grains in the intergalacticmedium,we will perform our analysis assuming a spherical geometry forthe grains.

Here we report on the analysis of a dedicated deep expo-sure withChandra on the bright, distant (z = 4.30) quasar QSO1508þ5714 (Moran & Helfand 1997 and references therein),designed to set constraints on the density of intergalactic dust.

2. DATA ANALYSIS

QSO1508þ5714was observedwithChandra on 2001 June 10for 88,971 s (effective exposure time), with the ACIS-S3 chip atthe focus. The data were processed with CIAO 2.2.1. We repro-cessed the data with the latest ACIS-S3 gain map, retaining onlyevent grades 0, 2, 3, 4, and 6. The backgroundwas stable through-out our observation, and we retained the entire exposure for anal-ysis. A close examination of the image of the source shows that itis definitely extended (Fig. 1). As is apparent from Figure 1, how-ever, the extended feature contains less than 10% of the sourcecounts and does not affect the radial profile of the source beyond"400, so it will not affect our measurement of a scattering halo.This jet, the highest-redshift X-ray jet known, is discussed else-where (Siemiginowska et al. 2003). In order tominimize the back-ground, we excluded events outside the energy range from 300 to8000 eV (the instrumental background in the S3 chip rises steeplyoutside these boundaries). We used the CIAO tasks celldetectand dmfilth to find and eliminate point sources in the image;only five weak sources were found within 30 of the quasar posi-tion. We measured the surface brightness in the image in a seriesof concentric annular regions centered on the peak brightness ofthe image (which coincides with the optical position of the qua-sar to within 0B5 in right ascension and declination; Hook et al.1995).

Our attempt to detect a faint extended scattering halo obviouslyrequires a careful treatment of the background. We retrieved thestandard ACIS-S3 quiescent background data (acis7sD2000-12-01bkgrndN0001.fits, 330 ks exposure) and reprocessed it with thesame gain map we used for the quasar image.We restricted eventenergies to the range 300–8000 eV and reprojected the back-

ground events using the aspect solution provided by the ChandraX-Ray Center for our observation. As a test, we subtracted asmoothed version of the resulting background image, simplyscaled by exposure time, from the quasar image, and found thatthe residual intensities were generally of order 1% of the back-ground image, or less. Avery faint diffuse structure (roughly cir-cular shape, diameter"500) remains, located at position angle 70#

(measured east fromnorth) at 2500 from the quasar, with a peak sur-face brightness of 0.2 counts arcsec$2. The total flux in this struc-ture is less than 0.1% of the quasar flux, so we did not subtract itfrom the image. Otherwise, the smoothed, background-subtractedimage appears completely dark outside a circle 2000 in radius cen-tered on the quasar. The background is slightly spatially inhomo-geneous (variations up to 10% from the mean surface brightnessoccur near the position of the quasar, on scales of "10), so wechose to use the same set of annular extraction regions as we usedfor the quasar to extract a background histogram. This histogramwas subtracted from the radial intensity profile of the quasar. Theresult is shown in Figure 2. We have not made corrections for theangular dependence of the telescope throughput, since these re-main small (a few percent) over the angular range of interest.The wings of the Chandra point-source response are domi-

nated by scattering due to microroughness on the mirrors. Theamplitude and angular distribution of the mirror-scattered lightdepend strongly on photon energy, and we therefore need to care-fully model the wings in the image of the quasar. In order to inves-tigate the intensity distribution,we compared the radial profilewiththe image of a low-redshift extragalactic point source with a spec-trum similar to QSO 1508þ5714.We selected the ACIS-S3 imageof 3C 273 (ObsID 1712); at redshift z = 0.158, this source is toonearby to have ameasurable intergalactic dust halo. ItsGalactic in-terstellar column density is small (NH % 1.7 ; 1020 atoms cm$2),so it does not have a Galactic dust scattering halo brighter than afew percent of the point-source flux, either; moreover, the Galac-tic scattering halo will be very extended and have very low sur-face brightness.We followed the same general procedures for theextraction of a radial intensity profile for 3C 273 as we used forQSO1508þ5714. There are no bright point sources near 3C 273,

Fig. 1.—ACIS-S3 image of QSO 1508þ5714 in the 0.3–8 keV band, binnedin 0B5 pixels. North is up, and east is to the left. The color scaling is logarithmic,with the brightest pixel containing 1370 photons, and the dimmest nonblackpixels containing one photon each. The position of the brightest pixel is atR.A. = 15h10m02.s90, decl. = þ57#02043B3 (J2000). The image is clearly extendedtoward the southwest. [See the electronic edition of the Journal for a color versionof this figure.]

Fig. 2.—Azimuthally averaged, background-subtracted 0.3–8 keV intensityprofile of the quasar, in counts divided by surface area (area measured in pixels0B5 on a side), vs. radial coordinate ( points with error bars). Superposed are aray-trace simulation for the core of the image (dotted line) and a radial profile forthe bright source 3C 273 ( jagged solid line); the downturn at r < 5 pixels (2B5) isa consequence of photon pileup in the detector. The upper solid curve representsa simple dust scattering halo model corresponding to an average dust densityof !d = 2 ; 10$6, for dust particle size 1 !m. The flat distribution extending outto approximately 300 pixels is a model halo for particle size 0.25 !m and!d = 2 ; 10$6.

PETRIC ET AL.42 Vol. 651

Petric et al. (2006)

over-idealized

d < 2 106

BUT

z = 4.3quasar

0.3 8 keVmost of the scattering

comes fromlower energy



The Astrophysical Journal, 751:93 (12pp), 2012 June 1 Corrales & Paerels

Table 4Observed and Expected Properties for QSO 1508+5714 and Its Scattering Halo

Energy N srcobs N sca

exp Nbkgc

(keV) (< 2′′)a (10′′–500′′)b

0.3–1.0 1943 375 (38)d 248091.0–2.0 1915 32 150742.0–3.0 637 3 122473.0–4.0 368 1 8479

Notes.a Number of counts observed within a 2′′ radius of point-source center.b The total number of scattered counts expected to be observed in the region10′′–500′′ away from the point source, rounded to the nearest integer.c The estimated number of background counts, extracted from a circular regionon the same image, for an area covered by an annulus 10′′–500′′ wide.d With a factor of 10 reduction due to absorption.

occurs for bright objects (!1 X-ray photon per second). Pile-upcauses an obvious depression of counts in the center of 3C 273,prevents the accurate measurement of source flux, and therebydegrades the accuracy of the radial profile within ∼10′′, or20 pixels (Gaetz et al. 2004). Since the radial profile from 3C 273appears to match QSO 1508+5714 best around 5′′ (10 pixels),the extrapolated wing profile is more likely to be an overestimateof the true PSF—perhaps explaining why the intensity from 3C273 matches the upper limits on the instrumental measurementbetween 30′′ and 60′′ (Petric et al. 2006, Figure 2).

We use the PSF calibration model by Gaetz et al. (2004) togauge the PSF + halo profile that would have resulted from thep = 4 dust distribution with density ΩIGM

dust = 10−5. A quickexamination of QSO 1508+5714 (ObsID 2241) using a circularregion with a 4 pixel radius yields approximately 5400 countsin the 0.3–8 keV range. (This is a lower limit to the total numberof source counts, as it does not consider counts spread intothe PSF.) For several energy bins, Table 4 shows the numberof point-source counts, expected number of halo counts in anannulus 10′′–500′′ centered on the point source, and the numberof background counts expected in that same region. The majorityof the scattered light will have photon energies of <2 keV.However, for soft X-rays (<1 keV) and large dust grains, theRG cross-section can overestimate the intensity of scatteringby a factor ∼10 due to absorption by grain material (Smith &Dwek 1998). For this reason, the 375 scattered counts expectedin the 10′′–500′′ annular region are likely too large by an order ofmagnitude. We therefore exclude photon energies below 1 keVfrom further evaluation. To take advantage of the soft X-rayband, the halo profile must be modeled with Mie scattering.

Figure 6 compares the radial brightness profile of QSO1508+5714 in the 1–8 keV energy band to the PSF anddust models. The magnitude of the error bars shows thatthe observations of Petric et al. (2006) are consistent withΩIGM

dust " 10−5 for the power-law distribution of dust grainschosen. However, the brightness profile is also consistent with apopulation of IGM dust composed purely of 1 µm sized grains.By including energies below 1 keV and using the RG cross-section, Petric et al. (2006) were expecting a much brighterX-ray halo. This was an overestimate leading to the limitΩIGM

dust < 2 × 10−6, which our analysis shows can be relaxed toΩIGM

dust # 10−5. This emphasizes the need for a more detailedmodeling, incorporating statistics and Mie scattering, to beimplemented in the future. An even more important task wouldbe the careful calibration of the Chandra HRMA PSF between1′′ and 20′′.

Figure 6. Top: the radial profile for QSO 1508+5714 in the 1–8 keV range,with the model point-spread function in gray. Bottom: the residual intensityafter subtracting the PSF model of Gaetz et al. (2004), compared to the intensityexpected from a p = 4 power-law distribution of gray dust (dashed line) and1 µm size grains (dotted line) with total mass density ΩIGM

dust = 10−5.

8. DISCUSSION

X-ray scattering provides a unique view of the cosmic historyof star formation, feedback, and the IGM that is complementaryto UV, optical, and infrared studies. A summary of our mainresults is as follows.

1. If dust is evenly distributed throughout the IGM, withΩIGM

dust = 10−5, then the dust-scattered X-ray light willbe ∼5% (E−2

keV) of the point-source brightness for objectsat z $ 2. The intensity of the scattered light is directlyproportional to ΩIGM

dust .2. Also, in the uniform IGM case, point sources at larger

values of z trade larger dust column densities for a narrowerhalo profile, making it slightly more difficult to distinguishscattered light from the point source.

3. For high-z point sources, scattered light comes predomi-nantly from intergalactic dust at z " 2. If the majority ofIGM dust came from star formation within sub-L* galaxiesat z > 3, then assuming a constant comoving density of dustgrains may still be reasonable for point sources at z = 4.

4. If X-ray light scatters through a dense, dusty clump inintergalactic space, then the resulting halo image will havea flatter radial profile compared to the uniformly distributedcase. The further away the clump is from the backgroundpoint-source, the larger the halo image will appear to anobserver.

5. If the dust mass to metallicity ratio is relatively constant, asobserved locally, then a typical DLA will have an opticaldepth to 1 keV X-ray scattering ∼1%–3%, proportional to(1 + zg)−2. This result scales with metallicity, so DLAs withnear-solar abundances may have τx ! 5%.

6. Dusty clumps at z ! 1 would have to be $200 kpc insize for a halo image to appear at angles larger than 10′′

from the point-source center. X-ray scattering thus offersthe opportunity to test if dusty outflows exist near quasarabsorption systems, but the outflow may require columndensities of dust that are so far unprecedented. Perhapsan X-ray point source whose line of sight is close to aforeground galaxy can test the ability of large grains to beexpelled from L* galaxies.

10


z = 4.3quasar

whereRG-Drude

approximationis valid

1 8 keV

can stillbe valid

(relax constraint)

d 105

Conclusion: Need to use MIE scattering solution




YES for a population of grey graphite grains:

AV 0.01 0.02

A1.78µm 0.01 0.02


2. Could we expect large (‘grey’) dust grains in the IGM?



M82 : Starburst galaxy imaged with Hubble, Spitzer, and Chandra



QSO

Ménard+ (2010)

Qua

sar c

olor

Magnification and dust reddening 7

Figure 4. Correlation between QSO reddening and galaxy overdensity as a function of angular scale. Note that the four independentcolors are taken from adjacent passbands and do not maximize the signal-to-noise ratio (see Figure 6 for such a quantity).

magnitude. No such behavior could be detected when us-ing stars instead of quasars as a control sample.

The amplitude of dust reddening depends only on theproperties and amount of dust around the foreground galax-ies selected for the cross-correlation. We can recover a similarwavelength-dependent extinction as in Figure 3 or the meanreddening signal presented in Figure 4 by using subsamplesof quasars in different magnitude ranges. As a further san-ity check, we also replaced our quasars with stars selected tohave the same magnitude and spatial distribution on the sky.Measuring the cross-correlation between star brightness andgalaxy overdensity with this sample, we find no appreciablecolor excess for galaxies with i < 20.5 and stars selectedin various magnitude ranges.3 Finally, we have investigatedthe dependence of the signal as a function of Galactic red-dening: by splitting the dataset into two regimes of Galacticextinction given by the Schlegel et al. (1998) map, we didnot detect any significant change in our signal.

4.1 Reddening

4.1.1 Scale dependence

We isolate the reddening effects by measuring the correla-tion between QSO color and foreground galaxy overdensity,

3 We have detected, however, some excess reddening for starswhen fainter galaxies (20.5 < i < 21) are used. This effect islikely to be attributed to a small contamination of faint galax-ies in the stellar sample at faint magnitudes where star-galaxyclassification is incomplete. This is expected to occur at i ∼ 21.The presence of galaxies in the star sample may give rise to asimilar correlation between source reddening and galaxy over-density, since galaxy clustering gives rise to an excess of reddergalaxies in overdense regions, which produces a similar signal thatmimics the reddening by dust. We have measured the reddening-clustering correlation of SDSS galaxies with i > 20.5 and foundthat a contamination of galaxies at about 10% could explain thereddening signal seen around stars.

wαβ(θ), where α and β indicate two different pass bands.We estimate the errors by computing the color covariancematrix from bootstrap resampling. Note that the errors oncolors are smaller than the errors on brightness changes.As Figure 4 shows, quasar colors and galaxy overdensitiesare positively correlated, i.e. quasars appear to be redderwhen closer to high concentrations of foreground galaxies.The reddening effects are detected for all color combina-tions, from θ ! 0.1′′ to about 2, corresponding to phys-ical scales ranging from 50 h−1kpc to about 10 h−1Mpc.The measurement probes galactic radii well beyond the sizeof galactic disks. The amplitude of the effect is stronger inbluer bands. We find that a background source whose lightpasses at around 20 h−1kpc from a foreground galaxy (se-lected with i < 21) will be, on average, redder by E(g− i) !E(B − V ) ! 0.01 mag.

4.1.2 Wavelength dependence

The five SDSS filters allow us to constrain the shape ofA(λ), the extinction curve of the dust associated with thegalaxies, through four independent colors. We measure thewαβ correlations for two angular bins: 0.14 < θ < 0.80 ar-cmin and 0.8 < θ < 8.0 arcmin, which correspond to ef-fective projected radii of about 20 < rp < 100 h−1kpc and100 h−1kpc < rp < 1h−1Mpc. In Figure 5 we show the cor-responding color excesses with respect to the r band. Wecompare these reddening measurements to the standard ex-tinction curve by fitting these data points with the func-tional form of the extinction curve provided by O’Donnell(1994). Such extinction curves are usually characterized bythe parameter RV = AV /E(B − V ), which characterizesthe slope of the extinction curve. The coefficient AV quan-tifies the amount of dust through its extinction in the Vband. The best fit for AV and RV is shown with the bluecurve. On small scales our measured reddening correspondsto AV = (1.3 ± 0.1) × 10−2 mag and RV = 11.3 ± 7.5, i.e.

c© 0000 RAS, MNRAS 000, 000–000

Angular distance from center of foreground galaxy

1 Mpc (z = 0.36)




YES for a population of grey graphite grains:

AV 0.01 0.02

A1.78µm 0.01 0.02


2. Could we expect large (‘grey’) dust grains in the IGM?

PERHAPS due to radiation pressure driven winds (feedback)

The dust grains that are efficiently ejected AND survive the process would likely be larger (Davies 1998, Ferrara 1991)

3. X-ray scattering has potential to find exotic dust in exotic places

diffuse IGM

galaxy halos absorption systemsQSO



The Astrophysical Journal, 754:116 (8pp), 2012 August 1 Menard & Fukugita

Figure 7. Cosmic mass density of dust contained in Mg ii absorbers (green diamonds) compared with various other estimates of the dust density, contained in galacticdisks (Fukugita & Peebles 2004; Driver et al. 2007), in virial radii of galactic halos (Menard et al.), in intergalactic space beyond virial radii, and the total amountproduced in cosmic history (Fukugita 2011). The dashed curve refers to the estimate from a global fit using a parameterized form for the redshift dependence for theentire sample (see the text).

We quote as a representative value at a low redshift (z ≈ 0.5)

ΩMg iidust " 2.3 ± 0.2 × 10−6. (13)

As mentioned above, this provides us with a lower limit onΩhalo

dust as our analysis does not include contributions from linesof sight that (1) produce weak or no Mg ii absorption, and(2) intercept high column density clouds with the amount ofdust that obscures the background source. We saw that dustmissed for the former reason would increase the amount byapproximately 20%.

Figure 7 includes estimates of the amount of dust in theliterature. MSFR give the abundance of dust within the virialradius of representative galaxies at z ≈ 0.34 (Ωhalo

dust " 2.1 ×10−6). Fukugita (2011) estimated the dust-mass density withinthe virial radius integrating over all galaxies assuming a typicalluminosity function and that the dust amount produced isproportional to luminosity, 2.6 × 10−6 at the median redshiftof the sample they used, and also extended it to include thecomponent outside the virial radius, which is indicated byMSFR. The latter gives ΩIGM

dust ≈ 5 × 10−6, showing that thetotal abundance of dust is consistent with the amount ought tobe produced in star-forming activity in cosmic time, when addedto dust in disks Ωdisk

dust ≈ 3–4 × 10−6. It is noteworthy that thedust amount in Mg ii clouds is close to that within the virialradius of galaxies, suggesting that dust is predominantly held inMg ii clouds within the virial radius. This is approximately halfthe total amount of dust expected to be spread outside galaxies.We note that the calculation of both Ωhalo

dust and ΩMg iidust adopts

SMC-type dust and the same reddening-to-dust mass ratio.The estimate of Fukugita & Peebles (2004) refers to dust

residing in disks of local galaxies, and Driver et al. (2007)estimate the dust abundance in disk galaxies at z ≈ 0, bothindicated in the figure for comparison. This shows that theamount of dust in Mg ii absorbers is similar to that remaining ingalactic disks: much is blown out of galaxies.

5. THE HYDROGEN MASS DENSITYIN Mg ii ABSORBERS

The same formalism also applies to estimate the mass densityof neutral hydrogen associated with Mg ii absorbers (Lanzettaet al. 1995). We have

ρMg iiH i (z) =

(dN

dz

)

Mg ii

NH i mH

dX/dz. (14)

The neutral hydrogen column density of Mg ii absorbers hasbeen studied by Rao et al. (2006), who compiled about 200Lyman-α measurements of Mg ii absorbers. Using this sample,Menard & Chelouche (2009) showed that the median NH i ofMg ii absorbers is described by

NH i(W0) = (2.45 ± 0.38) × 1019(

W0

Å

)2.08±0.24

cm−2. (15)

Using this relation10 and Equations (12) and (14), we find

ΩMg iiH i " (1.5 ± 0.3) × 10−4 (16)

for the total sample which has median redshift z ≈ 1.11 TheH i abundance in Mg ii clouds may slightly decrease toward lowredshift, but the change is within the error. We thus compute theglobal dust-to-H i ratio for Mg ii absorbers,

ΩMg iidust

ΩMg iiH i

" 151 ± 15

. (17)

Including the mass contribution from Helium and heavierelements, it gives a dust-to-gas mass ratio of about 1/(70 ± 20),

10 We remark that this gives a column density of W0 < 2 Å clouds below theempirical threshold of star formation in galaxies derived by Kennicutt (1998).11 After the completion of this work, we became aware that Kacprzak &Churchill (2011) derived a similar (and consistent) estimate of ΩH i traced byMg ii absorbers.

6

Ménard & Fukugita (2012)



It is still a BIG challenge

(MIE scattering)




Cygnus X-3:




Future prospects



Dust in the wind 1067

Figure 7. Comparison of reddening maps in the three dust models, eachnormalized to the MSFR result at 1 h−1 Mpc. The regions shown are aquadrant of the simulation cube, 25 × 25 × 50 (h−1 Mpc)3. Red/blue regionsindicate heavily/lightly reddened fields through the simulation, with thecolour scale running from E(g − i) = 10−5 up to 0.1, logarithmically.

MSFR argue that the dust in Large Magellanic Cloud (LMC)-likedwarfs is insufficient to explain the magnitude of their reddeningsignal. We concur with this conclusion. Reproducing the MSFRdata in the hybrid model with a physical dust-to-metal mass ratiorequires including galaxies up to several times 1010 M#, far largerthan the ∼3 × 109 M# baryonic mass of the LMC (van der Marelet al. 2002). Furthermore, the No-Wind simulation predicts a galaxybaryonic mass function that is inconsistent with observations, withan excessive global fraction of baryons converted to stars (Oppen-heimer et al. 2010).6 We do not consider the hybrid dust model to benearly as plausible an explanation of the MSFR results as the Windmodel; we present it as a foil to illustrate what would be requiredto explain MSFR’s findings with dust in low-mass galaxies. For theWind model, the metals in low-mass galaxies contribute much lessreddening than the intergalactic metals (Fig. 5).

4 D ISCUSSION

For the Wind model to succeed, we require that the dust-to-metalmass ratio in the IGM be comparable to that in the ISM, allow-ing only ∼50 per cent of the ISM dust to be destroyed duringits expulsion from galaxies and subsequent residence in the IGM.The validity of this assumption is by no means obvious, as the de-struction time-scales for 0.01 µm dust grains by thermal sputteringare ∼107.5(nH/10−3 cm−3)−1 yr at T = 106 K (Draine & Salpeter1979, fig. 7), while wind particles in the simulation typically re-main in the IGM for ∼109 yr before re-accreting on to galaxies(Oppenheimer et al. 2010, fig. 2). However, the sputtering rates de-cline rapidly towards lower temperatures (e.g. a factor of 300 lowerat T = 105 K), and with the wind implementation used in this simu-lation most ejected gas never rises above a few ×104 K. Ultraviolet(UV) or X-ray background photons are another possible destruc-tion mechanism for IGM dust, but the intergalactic radiation fieldis much lower intensity than the radiation field dust grains alreadyencounter in galactic star-forming regions.

A detailed consideration of dust survival in the IGM is beyondthe scope of this initial investigation, but the MSFR results clearlyraise it as an important subject for further study. The combinationof their measurements with our models gives a fairly clear idea ofwhat is required: survival of a substantial fraction of ejected dust,and an extinction curve that has roughly the colour dependence ofISM dust. The temperature sensitivity of thermal sputtering couldlead to preferential destruction of ejected dust in the higher masshaloes that host a shock heated gas halo (see Birnboim & Dekel2003; Keres et al. 2005, 2009a; Dekel & Birnboim 2006). In theWind model, most wind particles in haloes with M < 1013 M#have T < 105 K, but about 2/3 of the wind particles in haloes withM > 1013 M# have T > 3 × 106 K. If sputtering does destroyintergalactic dust at these temperatures, it could produce distinctivedrops in the galaxy–reddening correlation when it is evaluated formassive galaxies or for galaxies in dense environments. The recentstudy of McGee & Balogh (2010), which examines the correlationof background quasar colours with projected separation from galaxygroups of varying richness, provides some hint of such an effect,but their innermost point is at r = 1 h−1 Mpc, close to the virialradius of typical group mass haloes. Moreover, the Chelouche et al.(2007) measurements provide direct evidence for dust survival in

6 The Wind simulation predictions are reasonably consistent with the ob-served mass function for galaxies with L < L∗ , though it still predictsexcessive galaxy masses above L∗ .

C© 2010 The Authors, MNRAS 412, 1059–1069Monthly Notices of the Royal Astronomical Society C© 2010 RAS

Zu+ (2010)

PSF

Chandrabackground

Diffuse IGM

Foreground galaxy

The Astrophysical Journal, 737:26 (14pp), 2011 August 10 Novak, Ostriker, & Ciotti

Figure 6. Eddington ratio as a function of time, for three different time intervals in the A2 simulation.(A color version of this figure is available in the online journal.)

Figure 7. Power spectrum of the dimensionless mass accretion rate MBH/MEdd for A2. The units of the y-axis are arbitrary. For frequencies higher than the inverseof several times the central accretion disk timescale, the power spectrum is proportional to 1/f, the frequency response of the low-pass filter applied by our centralaccretion disk. For lower frequencies, the power spectra are determined by the physics of gas input, cooling, and feedback. At these low frequencies, the powerspectrum is nearly flat: ∝ f −1/4. That is, the accretion onto the SMBH has a power spectrum nearly the same as white noise.(A color version of this figure is available in the online journal.)

gas surrounding the SMBH. This would imply that isotropicfeedback would result in the least SMBH growth.

Figure 9 shows final SMBH masses versus wind openingangle. The largest SMBH growth occurs when the wind is

10

Novak+ (2011)



Dust in the wind 1067

Figure 7. Comparison of reddening maps in the three dust models, eachnormalized to the MSFR result at 1 h−1 Mpc. The regions shown are aquadrant of the simulation cube, 25 × 25 × 50 (h−1 Mpc)3. Red/blue regionsindicate heavily/lightly reddened fields through the simulation, with thecolour scale running from E(g − i) = 10−5 up to 0.1, logarithmically.

MSFR argue that the dust in Large Magellanic Cloud (LMC)-likedwarfs is insufficient to explain the magnitude of their reddeningsignal. We concur with this conclusion. Reproducing the MSFRdata in the hybrid model with a physical dust-to-metal mass ratiorequires including galaxies up to several times 1010 M#, far largerthan the ∼3 × 109 M# baryonic mass of the LMC (van der Marelet al. 2002). Furthermore, the No-Wind simulation predicts a galaxybaryonic mass function that is inconsistent with observations, withan excessive global fraction of baryons converted to stars (Oppen-heimer et al. 2010).6 We do not consider the hybrid dust model to benearly as plausible an explanation of the MSFR results as the Windmodel; we present it as a foil to illustrate what would be requiredto explain MSFR’s findings with dust in low-mass galaxies. For theWind model, the metals in low-mass galaxies contribute much lessreddening than the intergalactic metals (Fig. 5).

4 D ISCUSSION

For the Wind model to succeed, we require that the dust-to-metalmass ratio in the IGM be comparable to that in the ISM, allow-ing only ∼50 per cent of the ISM dust to be destroyed duringits expulsion from galaxies and subsequent residence in the IGM.The validity of this assumption is by no means obvious, as the de-struction time-scales for 0.01 µm dust grains by thermal sputteringare ∼107.5(nH/10−3 cm−3)−1 yr at T = 106 K (Draine & Salpeter1979, fig. 7), while wind particles in the simulation typically re-main in the IGM for ∼109 yr before re-accreting on to galaxies(Oppenheimer et al. 2010, fig. 2). However, the sputtering rates de-cline rapidly towards lower temperatures (e.g. a factor of 300 lowerat T = 105 K), and with the wind implementation used in this simu-lation most ejected gas never rises above a few ×104 K. Ultraviolet(UV) or X-ray background photons are another possible destruc-tion mechanism for IGM dust, but the intergalactic radiation fieldis much lower intensity than the radiation field dust grains alreadyencounter in galactic star-forming regions.

A detailed consideration of dust survival in the IGM is beyondthe scope of this initial investigation, but the MSFR results clearlyraise it as an important subject for further study. The combinationof their measurements with our models gives a fairly clear idea ofwhat is required: survival of a substantial fraction of ejected dust,and an extinction curve that has roughly the colour dependence ofISM dust. The temperature sensitivity of thermal sputtering couldlead to preferential destruction of ejected dust in the higher masshaloes that host a shock heated gas halo (see Birnboim & Dekel2003; Keres et al. 2005, 2009a; Dekel & Birnboim 2006). In theWind model, most wind particles in haloes with M < 1013 M#have T < 105 K, but about 2/3 of the wind particles in haloes withM > 1013 M# have T > 3 × 106 K. If sputtering does destroyintergalactic dust at these temperatures, it could produce distinctivedrops in the galaxy–reddening correlation when it is evaluated formassive galaxies or for galaxies in dense environments. The recentstudy of McGee & Balogh (2010), which examines the correlationof background quasar colours with projected separation from galaxygroups of varying richness, provides some hint of such an effect,but their innermost point is at r = 1 h−1 Mpc, close to the virialradius of typical group mass haloes. Moreover, the Chelouche et al.(2007) measurements provide direct evidence for dust survival in

6 The Wind simulation predictions are reasonably consistent with the ob-served mass function for galaxies with L < L∗ , though it still predictsexcessive galaxy masses above L∗ .

C© 2010 The Authors, MNRAS 412, 1059–1069Monthly Notices of the Royal Astronomical Society C© 2010 RAS

Zu+ (2010)

The Astrophysical Journal, 737:26 (14pp), 2011 August 10 Novak, Ostriker, & Ciotti

Figure 6. Eddington ratio as a function of time, for three different time intervals in the A2 simulation.(A color version of this figure is available in the online journal.)

Figure 7. Power spectrum of the dimensionless mass accretion rate MBH/MEdd for A2. The units of the y-axis are arbitrary. For frequencies higher than the inverseof several times the central accretion disk timescale, the power spectrum is proportional to 1/f, the frequency response of the low-pass filter applied by our centralaccretion disk. For lower frequencies, the power spectra are determined by the physics of gas input, cooling, and feedback. At these low frequencies, the powerspectrum is nearly flat: ∝ f −1/4. That is, the accretion onto the SMBH has a power spectrum nearly the same as white noise.(A color version of this figure is available in the online journal.)

gas surrounding the SMBH. This would imply that isotropicfeedback would result in the least SMBH growth.

Figure 9 shows final SMBH masses versus wind openingangle. The largest SMBH growth occurs when the wind is

10

Novak+ (2011)

PSF

Chandrabackground

Diffuse IGM

Foreground galaxy



A systematic survey for intergalactic dust will affect:



cosmic dust and metal budget






theory of galaxy feedback






theory of galaxy evolution







magnitude and timescaleof AGN variability







high precision cosmology

magnitude and timescaleof AGN variability



Thank you for your attention.

Corrales & Paerels (2014)arXiv:1311.5588

Corrales & Paerels (2012)ApJ 751, 93

Peek, Ménard, & Corrales (2014)coming soon

[email protected]

http://astro.columbia.edu/~lia

mailto:[email protected]

mailto:[email protected]



Documents

X-ray Studies of Interstellar and Intergalactic Dustlia/Corrales_grad_colloquium.pdfX-ray scattering is a diagnostic tool for ISM structure ... D. Russeil: Star-forming complexes and