1. technology and computing

2015-3-18 Tenmon-Gakkai ASTRO-H session
ASTRO-H
Cluster Physics
Takayuki TAMURA
ISAS/JAXA
On behalf of
the ASTRO-H/SWG cluster task force
Galaxy Clusters : Facts
 
The largest virialized system in the
Universe.
 
R~ 1 Mpc, Mtot 1014-1015 Msun, Tdyn ~ 109 yr.
Galaxies, ICM, Dark Matter = 5:15:80 in
masses.
  Intracluster medium (ICM)=optically thin
plasma → X-ray emission, Te~107-108K
(~keV).
  X-ray lines from O, Ne…Fe, 1/3 solar
metallicity, ICM ← primordial+processed in
galaxies.
 
2015-03 T.Tamura
Goals of the ASTRO-H mission
(Takahashi+ 2012)
1. 
Revealing the large-scale structure of the Universe and its
evolution
 
 
2. 
3. 
4. 
observe clusters of galaxies, to reveal the interplay between the
thermal energy of the intracluster medium and the kinetic energy
of sub-clusters, from which clusters form; measure the nonthermal energy and chemical composition; and to directly trace the
dynamic evolution of clusters of galaxies.
distant supermassive black holes … Understanding the extreme conditions in the Universe Exploring the diverse phenomena of the non-thermal
Universe Elucidating dark matter and dark energy
 
map the distribution of dark matter in clusters of galaxies and
determine the total mass of galaxy clusters at different distances
(and thus at different ages), and will study the role of dark matter
and dark energy in the evolution of these systems. Doppler Measurements of Gas Motions
(Bulk/streaming & Turbulent)
Cluster Turbulence
 
 
 
5
Clusters form via infalls,
accretions, and mergers. →
heat the ICM via shocks.
The heating depends on the
ICM micro physics.
Develop gas turbulence, and
accelerate cosmic rays, →
diffuse radio and X-ray halos.
Cluster mergers → The
most energetic
system after the Big
Bang.
1 Mpc
Fig. 2. The large scale velocity field on a thin slice though the center of cluster SB
shown overtop the logarithm of gas density (image, contours). The maximum velocity
vector is 2090 km/s. The image is 6.4 Mpc on a side.
Simulations, Norman & Bryan 1999 vation
(Navarro, Frenk & White 1995; Bryan & Norman 1998a) that the mean
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T.Tamura
cluster temperatures were, on average, about 0.8 of its virial value.
Thus we see that the gas has not completely virialized and sizable bulk
motions exist. Since the mean entropy profile increases with increasing radius,
X-ray imaging -> gas dynamics
Shock/Cold front
Assumptions:
(1) adiabatic flow
(2) Velocity in the plane of the
sky　→　rare case
Requires high spatial resolution
Shock in 1E~0657-56, M~3±0.4
(Markevitch et al. 2002).
Cold front in A3667, M~1±0.2
(Vikhlinin et al. 2001).
Shoc
k
Cold
Front
Radius
(Arcsec.)
5
A2256, X-ray bright, double peaked
merging clustergalaxy radial velocities
Berrington et al. 2002 SUB Peak
SUB
MAIN
Cold front ?
ΔV 〜 2000 km/s
MAIN Peak
Chandra X-ray Contour + galaxy image
(Sun+ 2002). See also Briel+ 1991 2015-03 T.Tamura
Radio: a diffuse halo, relics, and tailed
emission from galaxies → dynamically
young (e.g. Rottgering+ 1994)
Discovery of Gas Bulk motion in
A2256 with Suzaku
(Tamura+ 2011)
Fe K line
Main
Sub
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X-ray line center energies 2015-03
differs
slightly but significantly
ASTRO-H SXS simulation
the two components in A2256
SXS 100 ks
Note the Suzaku-observed
energy shift is about 30
eV ~ 1500 km/s.
Using the SXS with an
energy resolution better
than 7eV, → measure gas
bulk motions in a fair
number of X-ray bright
clusters.
2015-03 T.Tamura
Gas and Galaxy radial velocities
pa_vh_r30_v3.qdp
0.015
Redshift
0.02
0.025
R<30’
A
50.500
50.000
D
NNE1
F_2
f N1
49.000
49.500
0.01
H_2
h
W1_2
W1
E1
kp
C
B
SW1
S1_2
S1
e
c)
G_2
g
SE1
R
=3
0’
(6
60
0
4
9
18
37
74
149
298
594
50
100
150
200
PA (deg., N−>E)
250
300
350
Tamura+
201412:13
ttamura 11−Sep−2013
2015-03
CfA redshift
94T.Tamura
galaxies (Huchra+ 1995, 2012) bright galaxies
AH/SXI+HXI
X-­‐ray (Thermal) Turbulence
Black Hole/AGN: Jets and infla=ng bubbles, upli[s the gas, weak shocks Gas Thermal Energy (Ioniza=on, conduc=on, mixing)
Dumping. Dissipate
energy at Re =λv/ν　~ 1 (?)
Shock Hea=ng, Mach ? Mass ? Scale ?
Non-­‐thermal Radio (← B) & X-­‐ray (IC← CMB) Cosmic ray, γ ~ 104, Tcool < 108 yr , How to accelerate ? Scattering on
MHD waves
Magne=c field Turbulence, Spacial spectrum (λ1/3)? Viscosity (ν) ? Reorder & amplify
References:
Kitayama+ 2014,
Cattaneo+ 2009
AH/SXS
Gas Mo=ons (Infall, Collision/Merger, sloshing) Galaxy Mo=ons Dark Ma,er Gravity, move with gas ? Galaxies ?
The Perseus core
Fabian+ (2008), Fig.1:
Image of the Hα emission　140’’ x 150’’. HST
Chandra radial fraction
difference (Fabian+ 2011)
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(Not Scaled)
Current Limits on Turbulent velocity
CCD ΔE ~ 120 eV →
2300 km/s @ Fe-K
normalized counts s−1 keV−1
(b) CS5
0.1
0.01
1.4
ratio
1.2
1
0.8
0.6
6
6.5
Energy (keV)
7
Suzaku Limits (TT+ 2014;
also Ota+ 2007)
  Use Line width of Fe-K line.
  No additional width is
required, only upper limit of
the velocity
  Limited by CCD resolution
and its calibration.
  Limits > 1000 km/s ~ sound
speed.
XMM/RGS constrains on
Vturb at the core (e.g.
Sanders+ 2013).
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0
5
Energy (keV)
10
7.9
8
Energy (keV)
8.1
8.2
Detect and locate ICM turbulence
gure 49: Left: A 1 Msec SXS simulated spectrum of a dwarf galaxy. We incorporate a range of hypothetical sterile neutrino lines
ed dotted) at 4.0, 8.0, and 10.0 keV for ⌃dm = 100 M pc 2 and sin2 2✓ = 10 10 corresponding to the flux 0.04, 0.65, and 1.6
10 6 ph s 1 cm 2 , respectively, within the field-of-view of SXS. The adopted value of ⌃dm is typical of local dwarf spheroidals
d that of sin2 2✓ lies close to the current observational
No
di↵use galactic emission nor observable line broadening is
data andlimits.
folded
model
Perseus
simulated
spectrum
(wabsmodels
＊bapec)
sumed. The sum (black solid) includes
instrumental
and cosmic
background
taken from the ASTRO-H internal release
xs cxb+nxb 7ev 20110211 1Gs.pha). Right: Same as the left panel, except that the line at 8 keV is zoomed in.
normalized counts s−1 keV−1
10
Perseus, 100ks, 3’×3‘ → 11,000 cts in He-­‐like Fe w
here Eobs is the observed energy of the line.
Turbulent Velocity
1
.2
z
y
j
vturb
vturb
vturb
vturb
= 0 km/s
= 100 km/s
= 300 km/s
= 1000 km/s
x
k
0.1
he line-of-sight component of the turbulente velocity measured by line broadening can be a↵ected by calibraon errors in instrumental broadening
due to the line spread function and uncertainties in thermal broadening
o p
haracterized by the ion temperature.
To quantify their impacts, let us decompose the observed FWHM of a spectral line, assuming a Gaussian
ofile for each component, as
2
2
2
2
Wobs
= Winst
+ Wtherm
+ Wturb
+ ··· ,
(9)
here Winst , Wtherm , and
broadening, thermal6.7
broadening, and turbulent
6.4Wturb are the FWHMs
6.5 of instrumental6.6
oadening, respectively. Their nominal values Energy
are
(keV)
sekiya 6−Jun−2011 12:36
Winst ' 5 eV,
!1/2
p
kT ion
Wtherm =
8 ln 2
Eobs ,
mion c2
!1/2
! 1/2 ✓
kT ion
mion
Eobs ◆
,
= 4.9 eV
5 keV
56 mp
6.7keV
✓v ◆
p
turb
Wturb =
8 ln 2
Eobs ,
c
!✓
vturb
Eobs ◆
= 5.3 eV
,
100 km/s 6.7 keV
(10)
Gain errors (〜2.0 eV) dominate if counts >100 & Vturb < 300 km/s (11)
(12)
here T ion and mion are the temperature and the mass of the ion producing the line observed at Eobs , mp is the
abundance peaking at the radius of ∼50 kpc and declining both towards smaller or larger radii as shown in Fig. 1. The last functional
form is closest to the radial abundance profile derived from the deprojection analysis of the Chandra and XMM–Newton observations
(Schmidt, Fabian & Sanders 2002; Churazov et al. 2003) under the
assumption of a single-temperature optically thin plasma emission
model. The abundance ratios used are those of Anders & Grevesse
(1989). Using these data we calculated an optical! depth from the
centre of the cluster up to a radius of 1 Mpc, τ = n i σ0 dr , where
n i is the ion concentration and the cross-section for a given ion is
√
πhre c f
σ0 =
,
(3)
#E D
Resonant Scattering and Gas Motions
  Strong resonant transition
toward the core:
absorbed and re-emitted, τ >1.
  Change line profile in energy
and in space, depending on
(1) column density → liner scale
(2) velocity field →　gas
dynamics XMM line ratio (Ka/Kb
He-like Fe) → Lack of scattering
→ M > 0.5.
 ASTRO-H will resolve
resonance line within the Hetriplet → model-independent
spectroscopy for the 1st time.
where
#E D = E 0
"
= E0
$
2
2kTe
Vturb
+
Am p c2
c2
#1/2
2kTe
(1 + 1.4AM 2 )
2
Am p c
%1/2
.
(4)
In the above equations E 0 is the energy of a given line, A is the
atomic mass of the corresponding element, m p is the proton mass,
V turb is the characteristic turbulent velocity, M is the corresponding
Mach number, r e is the classical electron radius and f is the oscillator strength of a given atomic transition. The wavelengths and
absorption oscillator strengths are taken from the compilation of
Verner, Verner & Ferland
(1996).Limits
The ionization
equilibrium2004)
is that
XMM
(Churazov+
of Mazzotta et al. (1998). The set of Fe lines with an optical depth
larger than 0.2 (for M = 0) is given in Table 1. From this table it
is clear that (i) the 6.7-keV line of He-like iron is by far the most
optically thick line for the case of pure thermal broadening and (ii)
strong turbulence makes resonant scattering effects negligible for
all lines as emphasized by Gilfanov et al. (1987). For other types
of abundance profiles, the main result is the same; the 6.7-keV line
has an optical depth of the order of 3 and accounting for turbulence
reduces the optical depth to values smaller than 1.
Figure 2.
profile) wit
scattering.
redistribute
The res
approach.
line emiss
v1.3.0 res
counted fo
scattering
fact that w
which has
in Sazono
separate b
centre. Th
due to res
is shown i
line intens
ter outskir
without re
Fig. 3 for a
abundance
Gas motions in the core of the Perseus cluster
3120
31
I. Zhuravleva et al.
Figure 4. The 5–9 keV spectrum of the 30 arcsec to 2 arcmin annulus
centred on NGC 1275 and fitted with APEC and MEKAL models.
Figure 3. Ratio of the 6.7-keV radial brightness profiles without accounting
for resonant scattering to the profiles including the effect of resonant scattering. The three curves correspond to the three models of the abundance
profiles shown in Fig. 1. Dotted vertical lines show the two regions used for
spectra extraction.
similar for all three cases; the flux in the line is suppressed by a
factor of up to ∼2 within the inner 100-kpc region and is enhanced
by ∼10–20 per cent outside this region.
3 SPECTRA
Obviously the easiest way to reveal the effect of resonant scattering
is to derive the line ratios for the central region. The He-like iron K α
and K β lines are ideally suited for this purpose because both lines are
due to the same ion of iron. Guided by Fig. 3 we accumulated MOS
spectra for two annuli, 0.5–2 and 2–4 arcmin, centred on NGC 1275.
The inner 0.5-arcmin region was excluded in order to avoid possible
contamination from the NGC 1275 nucleus. The spectrum from 5 to
9 keV was fitted with the APEC (Smith et al. 2001) and MEKAL (Mewe,
Gronenschild & van den Oord 1985; Mewe, Lemen & van den Oord
1986; Kaastra 1992; Liedahl, Osterheld & Goldstein 1995) models
in XSPEC v.11.2 (Arnaud 1996) and is shown in Fig. 4. Temperature,
heavy metal abundances (with the abundance ratios of Anders &
Grevesse 1989) and redshift were free parameters of the models. We
can see that for the MEKAL model there is a clear excess on the left
wing of the He-like iron K β line, where the He-like nickel K α line(s)
is blended with the iron line. In order to remove this discrepancy
we have to either raise the nickel abundance above the standard Ni
to Fe ratio (Anders & Grevesse 1989) by a factor of ∼2 or assume
that the 6.7-keV complex is suppressed by resonant scattering, thus
causing peculiarities in the observed line ratio. On the other hand,
the most recent APEC v1.3.0 model provides an almost perfect fit to
the whole 5–9 keV portion of the spectrum. A very similar situation
is seen for the spectrum of the 2–4 arcmin annulus. In what follows
we assume that we can rely on the predictions of the newest APEC
model, which has a significantly richer set of lines in the region of
interest. The best-fitting temperature and abundance values for the
two annuli are T e = 4.29 ± 0.05 keV, Z = 0.501 ± 0.007 relative
to the solar abundance (Anders & Grevesse 1989) and T e = 5.18
"
C
± 0.06, Z = 0.446 ± 0.007. The values of temperature are higher
than we would infer from fitting a broader 0.5–9 keV spectral band,
which is expected for projected spectra, given the radial dependence
of the temperature in the Perseus cluster (see equation 2).
The best-fitting value of the abundance for the 5–9 keV spectrum
is, of course, dominated by the contribution of the 6.7-keV line. We
fixed all parameters (except abundance) at their best-fitting values
and recalculated the abundance, first ignoring part of the spectrum
containing the 6.7-keV complex and secondly ignoring the part containing 7.9-keV complex. Because only the 6.7-keV complex is affected by resonant scattering, the different values of abundance in
the two fits would indicate an important role for scattering. The ratio of abundances calculated this way is shown in Fig. 5 with two
crosses (for two annuli). Resonant scattering is expected to increase
the ratio of abundances well above unity for the inner annulus. However, the observed ratio is consistent with unity, indicating that any
effects from resonant scattering are small.
Finally, we have demonstrated that projection effects do not significantly affect the determination of the spectral parameters by using the outer 2–4 arcmin annulus as a background for the inner 0.5–
2 arcmin annulus. The resulting spectrum has a lower best-fitting
temperature ∼3.8 keV, as expected, but the line flux ratio again does
not show any obvious anomalies.
4 RO L E O F G A S M OT I O N S
The resonant K α He-like iron line contribution to the 6.7-keV (6.6–
6.8 keV) complex varies from 40 to 52 per cent for the gas temperature range from 3 to 5 keV. Therefore, suppression of the resonant line flux in the inner regions due to resonant scattering would
strongly affect the intensity of the whole complex. On the other
hand, the 6.9- and 7.9-keV complexes can be treated as effectively
optically thin. The good fit of the spectra with the APEC model of an
optically thin plasma with the solar mix of heavy elements can be
considered as an indication that the resonant scattering effects are
suppressed. As noted by Gilfanov et al. (1987) turbulent motions
of the gas may significantly reduce the optical depth of the lines.
The effect is especially strong for heavy elements, which have thermal velocities much smaller than the sound velocity of the gas. For
Energy
resolution
In FWHM
120 eV → 5 eV
2004 RAS, MNRAS 347, 29–35
2015-03 T.Tamura
Dickey J
Dupke R
Ezawa H
Foster A
128
Gastalde
Gilfanov
13, 3
Kahn S.
Vol. 3
San F
Kalberla
Morr
Lodders
Loewens
Mazzotta
403
Molendi
J., M
Sanders J
Sanders J
Sanders
349,
Sanders
1797
Sazonov
191
SXS simulation (Zhuravleva+ 2013)
Figure 9. Simulated spectra from the core (0.5–1.5 arcmin annulus) of the
Perseus Cluster as Astro-H will obtain from 100 ks pointing observation.
We used the whole field of view (2.85 arcmin on a side) with excluded
Contribution to cosmology(WP:Ota & Bautz)
 
Independent cosmology probes:
① 
② 
 
 
 
 
Growth rate of cosmic density inhomogenities (cluster mass function) & the
cosmic expansion history. “growth-of-structure”, e.g. Vikhlinin+ 2009.
Cluster baryon fractions (angular diameter distance) & expansion history. E.g.
Sasaki 1996.
Accurate Mtot of many systems, ← mass-observable relations
and X-ray observations → Future surveys (eROSITA)
Uncertainty in H.E. assumptions limits the power in
cosmology
Simulations → Mach (bulk/turbulence) ~ 0.1-0.2 at 1/5-1/2
Rvirial. (e.g. Lau+2009). → <10-15% error in Mtot.
X-ray H.E. vs. G-lens masses → deviation from H.E. e.g
Okabe+ 2014 (recent Suzaku/Subaru WL).
Reduce systematic error on the total mass estimates
amount
ICM hydrostatic equilibrium
(H.E)? → reduce error on the
Mtot
→ precise cluster
cosmology
✓
◆2 ✓ µ ◆
p
v
turb
amount
where v2turb is
⇢gas kT/µmp is
ptherm
' 1.3 ⇥ 10
2
turb
100 km s
1
0.6
kT
5 keV
!
✓
◆2 ✓ µ ◆ kT ! 1
p
v
the one-dimensional
velocity dispersion
of turbulence,
⇢gas is
turb
turb
' 1.3 ⇥ 10 2
.
1
0.6 5weight,
keV
km smolecular
the thermalptherm
pressure, µ is the100
mean
and mp is
terms
contribute
the hydrostaticvelocity
mass via
where
v2turb is thetoone-dimensional
dispersion of turbulence, ⇢gas is the gas ma
⇢gas kT/µmp is the thermal pressure, µ is the mean molecular
and mp is the proton m
2 weight,
Ω
:
the
present
dark
r
d(p
therm + pturb )
X
terms contribute
to the hydrostatic massM(<
via r) =
.
energy density
G⇢gas
dr
2
r d(ptherm + pturb )
W0 ≡ pX/ρX a
M(< r) =
.
dr
constant
energyto measure forG⇢the
gas first time
Our primary
goal dark
is hence
the contribution of
equation
of state
support
with
agoal
percent
order
accuracy
to infer
itscontribution
impact onofthe
mass e
Our primary
is hence
to measure
for and
the first
time the
turbulence
(Vikhlinin
et al.
measurements
around
the2009)
bright cluster center will be limited by systemati
support with a percent
order
accuracy and to infer its impact on the mass estimates of
2015-03
T.Tamura
measurements
the brightbroadening
cluster centerfor
willwhich
be limited
by systematicestimation
errors on t
spread
functionaround
and thermal
a conservative
Current
Model＊bapec)
Perseus
modelTheoretical
spectrum (wabs
z
y
1
j
6.4
x
100 ks central 3’×3’ 5 eV resolu3on u
o
p
l
2015-03 T.Tamura
= 0 km/s
= 100 km/s
= 300 km/s
= 1000 km/s
t KLn
s
v
h
6.5
6.6
Energy (keV)
Sekiya & Tamura 20011
vturb
vturb
vturb
vturb
k
a
e
0.1
keV2 (Photons cm−2 s−1 keV−1)
10
w
6.7
*%&&+&*C&D
sekiya 6−Jun−2011
12:04
B;1 *%&&+&*C&#
9
30
8
28
7
26
24
6
22
5
20
4
18
3
14
1
12
0
1
2
3
4
10−6
10−5
10−4
10−3
0.01
0.1
0
normalized counts s−1 keV−1
High energy
processes
(WP:
Kawaharada &
Madejski)
16
2
1
2
5
Energy (keV)
10
5
6
7
8
10
9
Wide-band and sensitive
imaging spectroscopy
with SXS, SXI, and HXI.
Search for
- Very high temperature
gas (T>12 keV)
-  Inverse Compton
radiation by cosmic-rays
and magnetic fields.
20
Figure 32: Simulated observations of A3667 North West relic for a 200 ks exposure, including the non-thermal emission with = 2.0
as described in the text. Top left: The HXI field of view (90 ⇥ 90 ) which covers the NW relic and its edge (red box), overlaid on the
20-cm ATCA image (R¨ottgering et al. 1997). Top right: Simulated HXI raw image (NXB and CXB are included) of the potential sharp
edge in hard X-ray band. The image is binned to 20 ⇥ 20 pixels. The color scale is in counts per pixel and indicated in the legend on
the right. The PSF and vignetting e↵ects are taken into account. Bottom: Background (NXB+CXB) subtracted SXI (black cross) and
HXI (red cross) spectra. For reference, solid lines show the thermal spectra with kT = 5 keV.
ASTRO-H White Paper
Clusters of Galaxies and Related Science (Kitayama+ASTRO-H SWG)
Sec3on 3tle (coordinators)
Main targets & keywords
1. Nature of gas mo3ons in the X-­‐ray brightest galaxy cluster (Tamura, Werner, Allen)
Perseus core and offsets
2. A detailed view of AGN feedback (Simionescu & Matsishita) M87 core
3. Plasma kinema3cs and clusters masses (Ota & Bautz)
A2029, A2199, A1795+
4. Mapping gas flows and turbulence in merging galaxy clusters (Markevitch & Akamatsu)
Coma core, A3667, A754, A2319, A2256
5. High-­‐energy processes (Kawaharada & Madejski)
A3667 (relic & center), Coma, hard X-­‐ray, non-­‐thermal
6. Chemical composi=on and evolu=on (de Plaa & Sato)
Virgo, Centaurus, Perseus, NGC5044+
7. Detec=ng and characterizing the warm-­‐hot intracluster medium (Galeazzi & Kitayama)
Pair clusters in superclusters
8. A spectroscopic search for dark maCer (Kitayama, Tamura, Allen)
Perseus and clusters, Milky way, dwarf galaxies, right-­‐
hand neutrino
2015-03 T.Tamura