Observations of Dust in Space fromSpace
Takashi Onaka(University of Tokyo)
UV observations (extinction & depletion)
X-ray observations (scattering & absorption)
IR observations (emission & absorption)

Small particles produced by
Gas-evaporation method in the lab

Interstellar Extinction Curve
12

E(-V)/E(B-V)

10

-Rv

Tells us thedust size andcompositions

original schematics

220nm hump

8
6

Distinct feature
at 220nm (UV)4

FUV rise

2
0

~1/

-2
-4

0

2

4

6

1/ (m-1)

8

Carbonaceous
origin of the
220nm feature
Graphite?
10
PAH?

Mysteries of the 220nm humpConstant
peakwavelength
againstvarying width

Fitzpatrick& Massa(1986)
ApJ, 307, 286

No correlation between 0 and ~ 1/[{1/ – (/02)}2 + 2]

Fitzpatrick & Massa
1986 ApJ, 307, 286

Difficult to be interpreted by single component model2 components for 220nm hump?
Interplanetary Dust Particles(IDPs)GEMS (glass withembedded metal sulfides)
an interstellar dust analogue?

Bradley et al. (2005) Science 307, 244

220nm in IDPs
A, B: Observations
C, D: hydroxylated (OH-)
Mg silicates
E: organic C in IDP
F: GEMS in IDP
OH-silicate has a feature
at 220nm(Steel & Duley 1987, ApJ, 315, 337)

Contribution of 2 components(carbonaceous and silicate)
may account for the
non-correlation of o and Interstellar Gas Depletion
Observations of interstellar gas absorption
 gas abundance in the ISM[metal absorption lines are mostly in the UV]
Depletion = Log (gas abundance in the ISM/reference abundance)
Reference abundance = generally solar abundance
Depleted elements are thought to reside in dust grains
Elemental depletion indicates the dust composition(associated with the uncertainty
in the reference abundance)

UV absorption lines
Observations
towards HD93521
HST/GHRS

Need highspectralresolution toresolve velocity
components
Most metal lines
are in the UV
Savage & Sembach(1996) ARAA
34, 279

Interstellar Depletion*temperature atwhich a half ofthe element inquestion will becondensed

Savage & Sembach(1996) ARAA, 34, 279

Deplet ion = Log[ N/ Nsolar]

Depletion correlatedwith condensation
Temperature Td*

1
Te
Tl

0

Kr

S

N
O

Ar

Pb

C

-1

Se
Sn Cl
Zn
B

Ge
Na
Ga

As
P
K

Si
Mn

Cu

Mg

Li

-2

Cr
Co

-3

V
Fe
Ni
Ti
Ca

-4

0

400

800

Td (K)

1200

1600

Depletion with abundance
original schematics

1

S

Depletion

0

Ar

-1

CO
Mg
Si

Na
Mn

-2
-3

Major elements in interstellar
grains are C, O, Fe, Si, and Mg
-4

N

0

2

4

Cr
Ni

6

Fe

Ca

Log (abundance; H=12)810

X-ray observations of interstellar dust
X-ray halo (angular scale ~1’)

Forward scattering by interstellar dust
produces a halo around the X-ray sourcee.g. Hayakawa (1970) Prog. Theor. Phys. 43, 1224
Mauche & Gorenstein (1986) ApJ, 302, 371
Predehl & Klose (1996) A&A, 306, 283(2) X-ray spectroscopy of absorption edge
& X-ray Absorption Fine Structure (XAFS)(E ≤ 10eV)
Abundance estimate of elements in gas and solid
phases without the reference abundance uncertainty

X-ray Halo

Forward scattering dominanty
1-x

X-ray source

sca
obs
x Observersca ~ obs + y/(1-x),= [1/(1-x)] obs
Emax

a(max)

obs = y/x

1

f(x)
dσ
I(θobs )  FxNH  dE S(E)  da n(a)  dx
(θsca )
2
dΩ
(1

x)
Emin
a(min)
0

Fx: source flux, n(a): size distribution
f(x): spatial distribution;
: scattering cross-section;
S(E): source spectrum

Scattering cross-section
2

6

2

 2




d
Z
a
F
E


2
(
)




6
 cm
 
 
 1.110  
 exp 
2 
3  d M   3g cm   0.1m   Z 
 2 ' 
2

2

2

4

   a   1 keV   F ( E ) 
2
11  2 Z  cm  6.310   
3  

M
3
g
cm
0
.
1
m


 
  E   Z 
 
2

2

2

’ = 10.4 (1 keV/E) (0.1m/a) arcmin|m-1| <<1, 2a|m-1|/ << 1; Rayleigh-Gans approximation[valid for E>1-2KeV]
F(E): atomic scattering factor ~ Z far from abs. edge
Mauche & Gorenstein (1986) ApJ, 302, 371
sca ~ a4 2/E2 ~ M2/a2/E2 ( ∝ M/a3)

Larger dust dominates in X-ray halo scattering
X-ray halo decreases with E-2
Porous dust has less scattering per mass-> information on the dust porosity/sizeObservations by ROSAT(0.1-2.4keV)

Point source response

halo

Predehl & Schmitt (1995) A&A, 293, 889

Correlation with Av
Ih/(Ix+Ih) = [1-exp(-sca)]sca=0.057Av+0.09

Predehl & Schmitt 1995,
A&A, 293, 889

Correlation with NH
sca = 0.05 NH(1021cm-2) – 0.83
X-ray halo produced byinterstellar dust

Predehl & Schmitt 1995,
A&A, 293, 889

Observations of X-ray
Absorption edge
Dispersion relation for f1<->f2 (F = f1 + if2 ) f 2 ( )f1 ( E )  Z   2d2
 0 E 2

Absorption edge
L
K
C
0.28 keVO
0.53 keV
Mg ~ 0.05 keV
1.3 keVSi
0.1 keV
1.8 keVFe ~ 0.7 keV
7.1 keVAbsorption edge of Fe
Absorption edge

Fe

f2

f1

f1 becomes a dip

Henke 1982 Atomic Data andNuclear Data Tables 27, 1

Scattering cross-section decreases
at the absorption edge

Absorption Features in sca
FeMgSiO4

Mg K
OK

original schematics

Fe L

Observations of Cyg-2
Chandraobservationsof Cyg X-2
aroundoxygenK-edgeCan beaccounted
for bya 3-edge
model
Takei et al.
(2002)
ApJ, 581, 307

Alternative interpretation

Spectrum can be interpreted by lines of O+ and O++(Juett et al. 2004, ApJ, 612, 308)

Need observational data of better statistics& laboratory data of compounds(Costantini et al. 2005, A&A, 444, 187)

X-ray halo of Cyg X-2
O

Fe

Mg

Si
3.8’
4.8’
5.8’

(Costantini et al. 2005, A&A, 444, 187)

Dips due to O, Mg, and Si edges are seen in the halo
spectrum, indicating the presence of O-contained dust

X-ray absorption fine structure(XAFS)
Absorption edge in solids has fine structures
depending on chemical bonds and compositions
Absorption edge energy of solid is different from gas
Distinguishable from gas absorption,
providing a new method to estimate dust composition,
chemical state, and abundance
X-ray Absorption Near-Edge Structures(XANES); E ≤ 10-20 eV
Extended X-ray Absorption Fine Structures(EXAFS); E ≥ 10-20 eV

Structures around edge

Draine (2003) ApJ, 598, 1026

Observations of XAFS
Silicates

GX13+1XAFS
Dust abundance&
diagnostics
of chemical
state of element
in solid phase

Si XAFS

Ueda et al. (2005) ApJ, 620, 274

What can we see in the infrared?
opticalinfraredInfrared emission comes from warm objects
Dust grains are heated by stellar radiation
and become warm in the ISM

Infrared observations of dust grains
Energy balance of dust grains∫(, a) Jd = ∫ (, a) B(T)d 
absorption=
radiation(, a): absorption cross-section of the grainJ： incident radiation
For interstellar radiation fieldJ ~  Wi F* (Ti),
Wi~10-13~10-14 :dilution factor
F*: stellar radiation, Ti = 7500, 4000, 3000K, +(Mathis et al. (1983) A&A, 128, 212)

T ~ 20K

Orion@140m by AKARI ®JAXA

http://www.ir.isas.jaxa.jp/ASTRO-F/Outreach/results/results.html#1119

Galaxy at optical

Credit & Copyright Axel Mellinger

Absorption by dust grains appears as dark lanes

Galaxy at 1.2m
original schematics

COBE@NASAIn near-infrared, most light comes from oldstars (red giants) with less extinction

Galaxy at 12m
original schematics

Diffuse emission from dust grains (~200300K) in our solar system dominates in mid-IR

Galaxy at 140moriginal schematics

Diffuse emission from interstellar dust (~20K)
dominates in FIR, well tracing ISM distribution

Mass estimate from FIR observations
FIR Emission: I() =  B(T) (optically thin case) = N Cabs = M k / D2,(for a << a: dust size, Cabs ∝ volume)
N: column density, M: total dust mass, D: distancek: mass absorption coefficient of dust (cm2/g)(= Cabs/volume/; : specific weight of dust)
k is independent of a, for a<<
From I() with k, T, & D, M can be estimated(typical k ~ 0.1 cm2/g @250m)
Hildebrand (1983) QJRAS, 24, 267

Be aware of the assumptions behind the formula:
optically thin, single temperature, k, …

Why do we observe in IR?
Origin of our Universe
Remote galaxies are red-shifted into IR
Search for extraterrestrial life
Planets are bright anda better contract against the Sun in IR
Less extinction in the IR and the emission from starforming regions peaks in the FIR (~100m)
ISM can be studied most efficiently in MIR to FIR,
including star-forming regions in our Galaxy & galaxies
IC1396@9+18m by AKARI

http://www.ir.isas.jaxa.jp/ASTRO-F/Outreach/results/results.html#1101

Material evolution in the ISM

Molecular Cloud
ISM
YSO
PN

SN->SNR
GiantsInfrared Observations from Space
Only a limited rangeof IR can be observed
from ground
Emission fromterrestrial
atmosphere disturbs
sensitive IR
observations
Cooled telescopes inspace provide ideal
facilities for IR
observations

from Astro-F booklet

Cooled Infrared Telescope in SpaceSpace mission has a stringent constraint on the weight
Launching conditions are not friendly for telescopes
Cryogenic distortion has to be carefully taken care ofLight-weight, strong, and small thermal distortion
mirrors are required for cooled telescope systems
Keeping a telescope cold on the orbit
requires new technology development(space cryogenics)
IC4954/4955@9+18m by AKARIInfrared Astronomical Satellite(IRAS)First IR satellite mission by USA, UK, & NL
60cm Be mirrors
Sun-synchronouspolar orbit
All-sky survey at
12, 25, 60, & 100m+ spectroscopy+ CPC
10 mo observation
with 500l LHe
http://coolcosmos.ipac.caltech.edu/image_galleries/IRAS/iras2.html@NASAAll-sky observations of IRAS

Observing the sky in the direction opposite toearth center; In 6 months, all-sky can be observed®D. IshiharaIR satellites after IRAS
MSX 1996(USA)
35cm AlCOBE/
DIRBE 1989
(NASA)
18cm Al
IRTS 1995(ISAS+ NASA)
15cm AlISO 1995(ESA+
ISAS+NASA)60cm fusedquartzSpitzer
2003(NASA)
85cm Be
AKARI 2006
(ISAS + ESA )70cm SiC
Survey; Observatory;
Survey + observatory

AKARI satellite
70cm SiC mirror
180L LHe + cryocoolers
enabled18 month cold mission(2006.2-2007.8)
JAXA mission withparticipation of ESA,
collaboration with UK, NL,& KoreaAll-sky survey surpasses
IRAS database+
Pointing observationsof imaging andspectroscopyin 2-180m

AKARI SiC light-weight mirror
Porous SiC core
CVD (Chemical Vapor Deposition) SiC coatweight~11kg

Vibration test

Ahhh

Cryo telescope test chamber
Cooling with
LHe
Interferometermeasuresdistortion
with coolingTelescope cryo-test

Kaneda et al. (2005) Appl. Opt. 44, 6823

Focal-plane instruments (FPIs)
Far-infrared surveyor Infrared Camera
(FIS)
(IRC)

FPIs

IRC
2–26 m

FIS
50–180 m

FIS

＠Nagoya Univ.

IRC

MIR-L

NIR

MIR-S

FIS

10 cm

Infrared Observations of interstellar
dust grains
尾中 & 石原 2009,

天文月報, 102, 103

AKARI 9/18m all-sky map

Diffuse Galactic Infrared Emission-4.5

25-60mexcess
6-12mUIR bands(polycyclic aromatichydrocarbons: PAH)

Log S(W/m2/sr)

>100mthermal emission-5

IRTS/MIRS

-5.5
-6
-6.5
-7

Stochastic heating 1
or IR fluorescence
model

COBE
Thermal Emission
10

100

Wavelength(m)1000

original schematics

Diffuse IR
emission spectraof theGalactic Plane

Surface Brightness (10-8 W/m2/sr/m)

Emission bands in MIR“Unidentified
Emission Bands”
at (3.3), 6.2, 7.7,
8.6, & 11.3m(+ more)Onaka et al. 1996
PASJ, 48, L59

1.4
b=0
b=1-2
b=2-3
b=3-5

1.2
1
0.8
0.6
0.4
0.2
0

4

5

6

7

8

9

10

Wavelength (m)

11

12

UIR band carriers(organic compounds in the ISM)
C-H stretch
3.3 m
C-C bend
6.2 m
C-C
7.7 m
C-H bend (in plane) 8.6 m
C-H bend (out of plane)
11.2 menhanced by ionization

Observed bands can be attributed to vibrations of
Polycyclic aromatic hydrocarbon (PAHs)
But exact nature of the carriers remains unclear
Relative band ratios should depend on physicalconditions (e.g., ionization) of the band carriers

Correlation of UIR bands with total
FIR intensity

Total FIR intensity is a product of the dust column
density and the incident radiation intensity
FIR =    B(T) dabsGʘxGo /4Go: radiation intensity in units of solar
neighborhood value Gʘ (= 1.6 x 10-6 Wm-2)
Stochastic heating or IR fluorescence model predicts
UIR band intensity is also a product of the column
density and the radiation intensity

UIR band strength   Go  FIR[Go ∝ T4+ abs100m = 700 -> Go]UIR7.7m/FIR in W51 by IRTS
1original schematics

Good correlation
with total FIR
0.1

Supporting
stochastic
heatingor fluorescence
model

0.01

0.001

1

10

Go

100

UIR bands in ionized regions

UIR7.7/FIR

ISO observationsof Carina & S171
regions

10-1

10-2

10-3

UIRs weaken
in ionized regions
Onaka et al. (2000)
ESA SP-456, 55

10-4 -4
10

Carina region
 Oph region
S171 region
10-3

10-2

[OIII]88m/FIR

10-1

PHT-S Results of Galactic Planel = ±5, ±15,±30, ±60& b = 0, ±1(49 quiet
regions)

Surface Brightness(10-6 W m-2 mm-1 sr-1)

No systematic variations in the UIR bands on GP

Kahanpää et al.
(2003) A&A 405, 999

Wavelength (m)IRAS100m & UIR7.7mArea I
Area II
Area IIIArea IVl ~ -10l ~ +50l ~ -130l ~ +170
IRAS100m Good correlation between UIR & 100m

4

2

0

-2

-4
-4

-6

-8

-10

-12 52

1000

2000

3000

4000 MJ y/ sr-7.00- 8.00- 9.00-10.00

50

48

500

UIR7.7m

1000

46

1500

44- 128

M J y/ sr

-130

0

-132

20

40

- 134

60

80

- 136176

MJ y/ sr

174

0

172

50

100

170

150

200

168

MJ y/ sr

3

2

1

0

-1

-2

-3

-5.00

-6.00

0.0e+00

1.0e- 06

2.0e- 06

W / m2/ sr-11.005250

48

0.0e+00 2.0e-07 4.0e- 07 6.0e-07 8.0e- 07

46

44 - 128

-130

- 132

-134

- 136

176

174

172

170

168

Sakon et al. (2004) ApJ, 609, 203
2.0e- 084.0e-086.0e-088.0e-08

5.00e-08

1.00e- 07

1.50e- 07

Surface Brightness (x10-7 W m-2 mm-1 sr-1)

Typical Spectra
30

Area I

20

Area II

20
10

10

Area IIIArea IV

2

1
1

0.5

4

6

8

10

12

4

6

Sakon et al. (2004) ApJ, 609, 203 Wavelength (m)

8

10

12

Variations in the UIR bands

UIR/UIR7.7

UIR/FIR

0.03
0.02
0.01

Area I
Area II
Area IIIArea IVUIR/FIRlarger inouterregion

0
1.2
1
0.8

11.3/7.7
larger inouterregion

0.6
0.4
0.2
0

4

5

6

7

8

9

10

11

12

Wavelength (m) Sakon et al. (2004) ApJ, 609, 203

UIR bands in galaxies
Emission from the Galactic plane shows tiny variations
External galaxies provide extreme conditions
Major UIR bands are at 3.3, 6.2, 7,7, 8.6, 11.2mubiquitously seen in spiral & starburst galaxies
The variations seen among normal galaxies are small
based on ISO observations
30

AKARI IRC
spectrum

25
20
15
10
5
0

original schematics
2

4

6

8

10

12

14

NGC1313@7m

UIR bands in elliptical galaxies
Elliptical galaxies are a matured, processed system
10

IRS spectrum

IRC spectrum NGC1316

NGC29741
5

10Wavelength (m)40

Kaneda et al. (2005) ApJL 632, L83

Kaneda (2007) ApJL, 666, L21

11.3 and 17m complex are clearly detected
6.2, 7.7, & 8.6m are very weakNo detection of 3.3m in NGC 1316 suggests
dominance of large & neutral PAHs

10

7.7m/11.3m band ratio variation instar-forming galaxies
Kaneda et al. (2008) ApJ 684 270

1

HII dominated

NGC1316

Low-luminosity AGN

0.1

Ellipticals

0.1

1[NeIII]/[NeII]

Smith et al. (2007) ApJ 656, 770
Kaneda et al. (2008) ApJ 684 270

Ellipticals (■) or
galaxies with AGN(▼) show lower
7.7/11.3m ratios
cf. AGN activity
in NGC1316 ceased
100Myr ago,suggesting effects
10
other than AGN also
play a role

Dwarf (young) galaxies (high [NeIII]/[NeII]) do notshow any systematic difference in the band ratio

AKARI observations of NGC6946
Interarm region
has a lower 7.7/11.3
ratio than arm region
Sakon et al. (2007)
PASJ, 59, S483

UIR band variations seen in galaxies(6.2, 7.7)/11.3m band ratio decreases
in inter arm regions
Elliptical galaxies associated with hot plasma(indicated by X-ray emission)
have extremely weak 6.2 & 7.7m bands
Ionization effects* or band carrier processing
in plasma environments?
•6.2 and 7.7m bands are enhanced•by ionization of PAHs
UIR bands in extended structures of galaxies(halo, jets, …) associated with ionized gas ?

M82 seen by AKARI
3.2m

7m

15m

4.1m

11m

24m

2kpc

Extended emission of dust ejected by superwind@NAOJ

original schematics

UIR bands in M82 filaments

Red: HGreen: 7m

3.2m (B)+ 7m(G) + 15m (R)Engelbracht et al. (2006) ApJL, 642, L127

UIR band detected in filamentary structures
seen in 7-15m extended over H features
PAHs survive in superwind?

UIR band ratio variation in M82
6.2, 7.7, & 8.6m
are well correlated
6.2 (7.7)/11.3m is low
outside the galactic plane
A similar trend seenin the interarm
regions of NGC6946
and the outer region of
our Galaxy
Galliano et al. (2008) ApJ 679, 310

NGC1569 seen by AKARI

Nearby starburst dwarf (12+log(O/H)=8.13)
Several filaments produced by galactic winds
3.2m
7m
15m

4.1m

300pc

11m

24m

Matsumoto et al. (2009) Proc. of IAU 251, 249

UIR band associated with filaments
H (G) + 7m (R)4m (B) + 7m (G) + 15m (R)Matsumoto et al. (2009) Proc. of IAU 251, 249

7m emission well correlated with Hfilaments
Band carriers present in filaments of ~106yr old,
where destruction should be very fast (~1000yr)
Produced from fragmentation in the shocked region?

AKARI Spectroscopy ofa filament 30of NGC1569
ON - OFF spectrum

25
20
15
10
5
1.2

0

1

2

4

6

8

10

Wavelength (m)

12

0.8

3.3, 6.2, 7.7, (8.6), & 11.3mbands + NIR excess aredetected in the H filament

0.6
0.4
0.2
0
2.5

3

3.5

4

4.5

5

5.5

Wavelength (m)Matsumoto et al. (2009) Proc. of IAU 251, 249

14

Spectra of Filament & Disk
30
25

0.4

IRC spectrumof the filament

IRS spectrum of
NGC1569 disk
0.3

20
15

0.2

10
0.1

5
0

6

7

8

9

10

11

12

0

5

6

7

8

9

10

11

Wavelength (m)
Wavelength (m)Matsumoto et al. (2009) Proc. of IAU 251, 249

12

13

7.7/11.3m ratio is smaller in the filament than in disk
7.7m band seems to be narrower

Comparison of band ratio
10

Disk
1

filament
HII dominated

NGC1316

Low-luminosity AGN

0.1

Ellipticals

0.1

1[NeIII]/[NeII]

10

Kaneda et al. (2008) ApJ 684 270

HII regions in M101

M101 HII regions provide
a large range ofionization & metallicity
Gordon et al. 2008 ApJ, 682, 336

UIR band equivalent width vsmetallicity & ionization

Metallicity

Gordon et al. 2008 ApJ, 682, 336

No clear correlation
with metallicity

Ionization

UIR band decreases
in highly ionized regions,indicating the decreaseof the band carriersNo correlation of band ratios

Gordon et al. 2008 ApJ, 682, 336

6.2/11.3 & 7.7/11.3 are expected to increase with theionization parameter,
but no systematic trend is seen in M101 HII regions
A trend of band carrier processing seen so far isonly the 6.2/11.3 & 7.7/11.3 decrease in hot plasma
environments (halo, jet, elliptical galaxies)

IR spectroscopy for the study ofgas abundance
Forbidden lines have negligible coupling with radiation
and are excited exclusively by collisions
Intensity is a function of local temperatureand density of the collision partner
Forbidden lines in IR
Less affected by extinction
E < kTe; insensitive to Te
Intensity is a function of density only
IR spectroscopy provides a unique opportunity to
study elemental abundance in densestar-forming regions
All Sky Map@9m by AKARILarge Si abundance in active regions(Mizutani et al. 2004, A&A,
423, 579)

S171 & Oph regions(Okada et al. 2003, A&A,
412, 199; Onaka et al. 2006,
A&A, 640, 383)

[SiII]35m intensity (Wm-2sr-1)[SiII]35m
ISO observations of
Carina region-6>10-25% Si in gas
phase even in highdensity regions(≤ ~5% in ISM)
Large amount of Si involatile grains otherthan silicates?
Then how is Fe?

2.5 102 10

Mizutani et al. 2004, A&A, 423, 579-6-6

1.5 10

1 10

-6

5 10

-7

50% Si relative to solar-3& ne < 50cm-3
0
0

5 10

-8

1 10

-7

-7

1.5 10

-7

2 10

-7

2.5 10[NII]122m intensity (Wm-2sr-1)

-7

3 10

Si & Fe gas abundance in HII regions
0.1

0.01

ISM depletion
0.001

0.1

Large gas
Si iscommonin HII
regions
Fe is ina different
component
fromvolatile Si

1
Okada et al. 2008, ApJ, 682, 416
+
++
Si /S

IR spectroscopy is efficient for the studyof gas abundance in dense active regions

Observations of a Supernova
with AKARISupernovae are potentially an important dust supplierPrevious observations indicate a significantly small
amount of dust formed in SNe compared totheoretical predictions
Dust formation in SNe is unclear at present
SN2006jc Type Ib
2006 October 9.75(UT) in UGC4904
AKARI observations were carried out with IRC~200 days after explosion (29 April 2007)
SNR B0104-72.3 in SMC @AKARI 4, 7, & 11m (Koo et al. 2007, PASJ, 59, S455)

AKARI observations of SN2006jc
3-7-11m color

Prism spectrum(2-5m)UGC4904
SN2006jc
Strong emission in MIR detected

I. Sakon et al. (2009) ApJ in press

Strong MIR emission fromcircumstellar dust
Best fit: Hot (~800K: SN dust) andwarm (~300K: circumstellar) carbon dust
Few (<10-4Mʘ)dust grains(hot component)
formedin SN 2006jc

Warm dust ~ 3x10-3Mʘ

Dust formationin SNe isinefficient?

I. Sakon et al. (2009) ApJ in press

Dust associated with SN remnants

SNe are a source of dust and also destroy IS dust
Observations of SNRs in our Galaxy are difficult
because of confusion -> Go to Large Magellanic Cloud

SNRs in the LMC by AKARI(Seok et al. 2008, PASJ, 60, S453)

LMC@AKARI 9, 18, 65, 90, &140m

Multi-component spectrum of SNRs
Tcold = 40K
Tcold = 30K

24 and 70m flux is well
fit by dust destruction
model of SNR shocks(Borkowski et al. 2006, ApJ, 642, L141;
Williams et al. 2006, ApJ, 652, L33)

15, 24, & 70m flux can
not be fit with a single
temperature gray body,suggesting a cold dust
component of large mass
Single T gray body

Seok et al. 2008 PASJ 60, S453

More dust associatedwith SNRs in LMC,
indicating less efficientdust destruction in SNRsIRS Spectra of SNRs
Cas A

Kepler

120
100
80
60
40
20
10

Rho et al. 2008, ApJ,. 673, 271

15

20

25

30

Wavelength (m)35

original schematics

SNR spectra have dust features + line emission
cannot be fit by a single temperature gray bodySpectroscopic study is important40

Dust processing in the ISM
Variations of the UIR bands provide significant
information on the ISM processingof carbonaceous grains
A systematic trend is seen only in the 6.2/11.3 and7.7/11.3 in particular environments
Role of SNe for dust supply and destructionis yet to be investigated observationallySpectroscopic observations are indispensable tocorrectly understand the IR emission from SNRs
IR line spectroscopy has a great potential to study
gas abundance of dense regions
Cyg X @AKARI 90 & 140m

CODEN:PASJA60(SP2) S375–S542 (2008)

ISSN:0004–6264

Special Issue: Recent Results from AKARI

PASJ, 60

http://www.akari2009.com

AKARI All-Sky Survey
9m source (~700,000)

9 & 18m

90m source (~60,000)

65, 90, 140,
& 160m

9-18m color map
http://www.ir.isas.jaxa.jp/ASTRO-F/Outreach/results/PR081119/pr081119_1.html

AKARI All-Sky Survey Performance
1000

10

IRAS
1

Pic at 100pc

100

AKARI
10

0.1

0.01

10

Wavelength (m)

100

1

10

Wavelength (m)

尾中 & 石原 2009, 天文月報, 102, 103Higher sensitivity in 9-18m than IRAS
Higher spatial resolution than IRAS
Point Source Catalogues will be released
to the public in 2009100

Herschel & Planck (ESA)(2009)

Herschel (3.5m) (55-672m)

Planck (1.5m) (350m – 9mm)

WISE & JWST (NASA)

WISE (2009) 40cm
3-24mJWST (2013) 6.5m
<28m

Space Infrared Telescope for
Cosmology and Astrophysics (SPICA)
SPICA 3.5m ISAS/JAXA
Cooled telescope bymechanical coolers
for wavelengths 5 -200m
Approved tophase-A study by JAXAESA contribution approved
aiming at 2017 launch
http://www.ir.isas.jaxa.jp/SPICA/index.html

Thank you for your attention
Many thanks to the LOC

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