The gas-to-extinction ratio and the gas distribution in ...We investigate the relation between the optical extinction (AV) and the hydrogen column density (N H) determined from X-ray

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Mon. Not. R. Astron. Soc. 000, 000–000 (0000) Printed 26 July 2017 (MN LATEX style file v2.2)

The gas-to-extinction ratio and the gas distribution in the Galaxy

Hui Zhu1,2,3, Wenwu Tian1,4,5⋆, Aigen Li3 and Mengfei Zhang1,41Key Laboratory of Optical Astronomy, National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China2Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA3Department of Physics and Astronomy, University of Missouri, Columbia, MO 65211, USA4School of Astronomy and Space Science, University of Chinese Academy of Sciences, Beijing 100049, China5Department of Physics and Astronomy, University of Calgary, Calgary, Alberta T2N 1N4, Canada

Accepted 000. Received 000; in original form 000

ABSTRACT

We investigate the relation between the optical extinction (AV ) and the hydrogen columndensity (NH) determined from X-ray observations of a large sample of Galactic sight-lines toward 35 supernova remnants, 6 planetary nebulae, and 70 X-ray binaries for whichNH was determined in the literature with solar abundances. We derive an average ratio of〈NH/AV 〉 = (2.08± 0.02)× 1021Hcm−2 mag−1 for the whole Galaxy. We find no corre-lation between 〈NH/AV 〉 and the number density of hydrogen, the distance away from theGalactic centre, and the distance above or below the Galactic plane. The 〈NH/AV 〉 ratio isgenerally invariant across the Galaxy, with 〈NH/AV 〉 = (2.04± 0.05)×1021Hcm−2 mag−1

for the 1st and 4th Galactic quadrants and 〈NH/AV 〉 = (2.09± 0.03)× 1021Hcm−2 mag−1

for the 2nd and 3rd Galactic quadrants. We also explore the distribution of hydrogen in theGalaxy by enlarging our sample with additional 74 supernova remnants for which both NH

and distances are known. We find that, between the Galactic radius of 2 kpc to 10 kpc, thevertical distribution of hydrogen can be roughly described by a Gaussian function with a scaleheight of h = 75.5± 12.4 pc and a mid-plane density of nH(0) = 1.11± 0.15 cm−3, corre-sponding to a total gas surface density of

∑gas ∼ 7.0M⊙ pc−2. We also compile NH from

19 supernova remnants and 29 X-ray binaries for which NH was determined with subsolarabundances. We obtain 〈NH/AV 〉 = (2.47± 0.04) × 1021Hcm−2 mag−1 which exceedsthat derived with solar abundances by ∼ 20%. We suggest that in future studies one may sim-ply scale NH derived from subsolar abundances by a factor of ∼ 1.2 when converting to NH

of solar abundances.

Key words: ISM: dust, extinction — ISM: supernova remnants, planetary nebulae — X-rays:X-ray binaries

1 INTRODUCTION

Interstellar extinction results from the absorption and scattering

of starlight by interstellar dust grains. The knowledge of the in-

terstellar extinction (Aλ) at wavelength λ per hydrogen column

(NH), particularly in the ultraviolet (UV), visible, and near in-

frared, is crucial for correcting for the interstellar obscuration and

restoring the true luminosity of astronomical objects, and for de-

termining the dust-to-gas mass ratio and the photoelectric heating

rates of interstellar gas. So far, the commonly adopted hydrogen-

to-extinction ratio is that of Bohlin et al. (1978), NH/E(B−V ) =5.8 × 1021 Hcm−2 mag−1, where E(B − V ) ≡ AB − AV ,

the colour excess, is the difference between the B band extinction

(AB) and the V band extinction (AV ). The colour excess E(B−V )

⋆ E-mail: [email protected]

relates to AV through RV ≡ AV /E(B−V ), the total-to-selective

extinction ratio.

Bohlin et al. (1978) derived the hydrogen-to-extinction ratio

from a sample of 100 stars with E(B− V ) up to ∼ 0.5mag based

on the ultraviolet (UV) absorption spectra of HI and H2 observed

by the Copernicus satellite. With RV ≈ 3.1 for the Milky Way

(MW) diffuse interstellar medium (ISM; Whittet 2003), the Bohlin

et al. (1978) hydrogen-to-extinction ratio becomes NH/AV ≈1.87× 1021 Hcm−2 mag−1. Similarly, as summarized in Table 1,

Jenkins & Savage (1974), Whittet (1981), Diplas & Savage (1994),

and Rachford et al. (2009) derived the hydrogen-to-extinction ra-

tio from the UV absorption spectra of HI and H2 measured by

the Orbiting Astronomical Observatory 2 (OAO-2; Jenkins & Sav-

age 1974), the Copernicus satellite (Whittet 1981), the Interna-

tional Ultraviolet Explorer (IUE; Diplas & Savage 1994), and the

Far Ultraviolet Space Explorer (FUSE; Rachford et al. 2009). The

NH/AV ratios have also been derived by Sturch (1969), Knapp &

c© 0000 RAS

http://arxiv.org/abs/1706.07109v2

2 Zhu, Tian, Li & Zhang

Table 1. NH/AV Ratios Determined from the HI and H2 UV Absorption Spectra.

Hydrogen NH/AV Sample Reference Comments

Type (1021 cm−2 mag−1 H) Size

HI+H2 2.15±0.14 17 Rachford et al. (2009) FUSE observations of stars with E(B − V )<∼0.85mag

HI 1.59±0.09 393 Diplas & Savage (1994) IUE observations of stars with E(B − V )<∼0.6mag

HI+H2 1.9 12 Whittet (1981) Copernicus observations of stars with E(B − V )<∼0.6mag

HI+H2 1.87 75 Bohlin et al. (1978) Copernicus observation of stars with E(B − V )<∼0.5mag

HI+H2 2.0 95 Jenkins & Savage (1974) OAO-2 observations of stars with E(B − V )<∼0.4mag

Table 2. NH/AV Ratios Determined from the HI 21 cm Emission and/or CO emission.

Hydrogen NH/AV Sample Reference Comments

Type (1021 cm−2 mag−1 H) Size

HI 2.68 – Liszt (2014) Galactic latitudes |b|>∼20◦ and E(B − V )<∼0.1mag

HI 1.56 76 Heiles (1976) stars and globular clusters with E(B − V )<∼0.6mag

HI 1.6 81 Knapp & Kerr (1974) globular clusters with E(B − V )<∼1.3mag

HI ∼ 1.3–1.6 40 Sturch (1969) RR Lyrae stars with E(B − V )<∼0.4mag

HI+H2 2.4 – Chen et al. (2015) Galactic latitudes |b|>∼10◦ and E(B − V )<∼1.0mag

Figure 1. Hydrogen column densities NH against the optical extinction

AV for the AG89 sample. Red circles: SNRs, Blue rhombus: PNe, Black

triangles: XBs.

Kerr (1974), Heiles (1976), Liszt (2014), and Chen et al. (2015)

based on the column densities of hydrogen measured by CO and

the 21 cm emission of HI.

While the aforementioned studies directly measured the col-

umn densities of HI and H2 (see Table 2), an alternative way of de-

termining NH is to measure the X-ray absorption spectra of heavy

elements such as O, Mg, Si and Fe. One derives the column densi-

ties of these heavy elements from the measured X-ray spectra and

then converts to NH by assuming an interstellar reference abun-

dance standard which specifies the relative abundances of these

heavy elements (i.e., O/H, Mg/H, Si/H, and Fe/H). This method

take ionized, neutral and molecular hydrogen into consideration

and has been applied to measure NH toward supernova remnants

(SNRs), X-ray binaries (XBs), star-forming (SF) regions, and ex-

tended sources (see Table 3).

However, the NH/AV ratios derived from the X-ray absorp-

tion data involves the unknown interstellar abundance X/H (where

X = O, Mg, Si, Fe ...), often expressed in units of ppm (i.e., parts per

million). One often assumes the interstellar elemental abundance

to be solar. However, the solar abundances of the key heavy ele-

ments such C, O, N, S, Mg, Si, and Fe reported in the literature

have undergone substantial changes over the past four decades (see

Li 2005, Asplund et al. 2009). Also, there are arguments that the

interstellar abundances could be better represented by those of B

stars (because of their young ages) which are just ∼ 60–70% of the

widely adopted solar values (“subsolar”; see Snow & Witt 1995,

1996, Sofia & Meyer 2001). As shown in Table 3, from the same

set of X-ray data one would derive a much higher (by > 30%)

NH/AV ratio if one adopts the reference abundance of Wilms et

al. (2000; hereafter W00) for O and Fe instead of that of Anders

& Grevesse (1989; hereafter AG89). Also, these previous studies

suffer from small sample sizes and/or limited space distributions of

the sample sources.

To reduce the above-mentioned limitation and to obtain a more

reliable NH/AV ratio, in this work we carefully compile NH and

AV from the literature by separating those with NH derived from

solar abundances (AG89) from those with NH derived from subso-

lar abundances (W00). We finally arrive at a sample of 35 SNRs, 6

planetary nebulae (PNe), 70 XBs for which NH was derived with

solar abundances (AG89) and 19 SNRs, 29 XBs for which NH was

derived with subsolar abundances (W00). We present our sample

in §2 and derive a renewed mean 〈NH/AV 〉 ratio in §3. Also in

§3 we estimate the gas-to-dust mass ratio and explore the spatial

variation of 〈NH/AV 〉 and the gas distribution in the Galaxy. The

major results are summarized in §4.

2 COMPILATION OF AV AND NH

2.1 AV

If a group of lines are from the transitions with the same up-

per level, their intensity ratios would only depend on the transi-

tion probabilities. Thus they can be employed to calculate the ex-

tinction toward extended objects, e.g., SNRs and PNe. The fre-

quently used line ratios include Hα(6563 A)/Hβ(4861 A), known

as the Balmer decrement (e.g., SNR G67.7+1.8, Mavromatakis

c© 0000 RAS, MNRAS 000, 000–000

The Galactic NH–AV Relation 3

Table 3. NH/AV Ratios Previously Determined from Interstellar X-Ray Absorption Data.

NH/AV Sample Interstellar Abundance O/H Mg/H Si/H Fe/H Reference Comments

(1021 H cm−2 mag−1) Size Standard (ppm) (ppm) (ppm) (ppm)

2.87±0.12 17 Wilms et al. (2000) 4.90 2.51 1.86 2.69 Foight et al. (2016) SNRs with E(B − V )<∼10 mag

2.08±0.30 23 Anders & Grevesse (1989) 8.51 3.80 3.55 4.68 Valencic & Smith (2015) XBs with E(B − V )<∼2.5 mag

2.2+0.3−0.4

>100 Anders & Grevesse (1989) 8.51 3.80 3.55 4.68 Watson (2011) γ-ray bursts

2.21±0.09 22 Anders & Grevesse (1989) 8.51 3.80 3.55 4.68 Guver & Ozel (2009) SNRs with E(B − V )<∼10 mag

∼ 1 20 Anders & Ebihara (1982) 7.39 3.95 3.68 3.31 Vuong et al. (2003) SF regions with distances<∼0.5 kpc

1.79±0.03 26 Anders & Ebihara (1982) 7.39 3.95 3.68 3.31 Predehl & Schmitt (1995) SNRs and XBs with E(B − V )<∼6.1 mag

2.19±0.52 5 - - - - - Ryter et al. (1975) SNRs with E(B − V )<∼10 mag

2.22±0.14 7 - - - - - Gorenstein (1975) SNRs with E(B − V )<∼10 mag

1.85 4 - - - - - Reina & Tarenghi (1973) XBs, extended sources with E(B − V )<∼10 mag

Figure 2. Left (a): NH–AV relation derived from the XBs of the AG89 sample. Right (b): Same as (a) but for SNRs.

Figure 3. Left (a): NH–AV relation derived for the whole AG89 sample but with the three outlying SNRs excluded. Right (b): Same as Figure 2b but with

the three outlying SNRs excluded.

c© 0000 RAS, MNRAS 000, 000–000


Table 4. Optical extinction AV and hydrogen column density NH compiled from the literature (see Appendix for details).

Object Gla Gba NH(KBH)b NH(DL)c NH(AG89)d NH(W00)e AV Df NH(AG89)/AV NH(W00)/AV

Name (degree) (degree) (1021 Hcm−2) (1021 H cm−2) (1021 Hcm−2) (1021 Hcm−2) ( mag) (kpc) (1021 H cm−2 mag−1) (1021 H cm−2 mag−1)

Supernova remnants

G4.5+6.8 north 4.5 6.8 2.12 2.38 5.1±0.3 5.4±0.1 3.0±0.4 5.1 1.89±0.27 1.93±0.28

G6.4-0.1 central 6.4 -0.1 11.40 14.40 4.3±0.1 6.3±1.9 3.3±0.3 1.9 1.30±0.12 1.91±0.60

G6.4-0.1 northeast 6.4 -0.1 11.40 14.40 7.7±0.2 4.0±0.5 2.0 1.93±0.25

G11.2-0.3 southeast 11.2 0.3 12.70 17.50 15.7±1.5 13.0±1.7 4.4 1.21±0.20

G13.3-1.3 13.3 -1.3 8.05 8.11 5.5±4.5 0.5±0.1 3.0 11.00±9.26

G39.7-2.0 east jet 39.7 -2.0 6.49 7.63 5.2±0.1 2.3±0.3 5.5 2.26±0.30

G53.6-2.2 central 53.6 -2.2 5.48 6.69 6.1±0.1 3.1±0.4 2.3 1.97±0.26

G53.6-2.2 southwest 53.6 -2.2 5.48 6.69 4.8±0.8 6.7±0.3 4.0±0.5 2.3 1.20±0.25 1.68±0.22

G54.1+0.3 54.1 0.3 11.90 14.30 18.1±0.5 25.5±0.4 7.3±0.1 6.2 2.48±0.08 3.49±0.07

G67.7+1.8 north 67.7 1.8 9.42 10.50 4.1±0.9 1.7±0.4 2.41±0.78

G69.0+2.7 central 69 2.7 6.59 7.92 3.0±0.1 4.5±0.2 2.5±0.3 1.5 1.20±0.15 1.80±0.23

G74.0-8.5 74.0 -8.5 1.60 1.68 0.34±0.03 0.25±0.06 0.54 1.36±0.25

G78.2+2.1 78.2 2.1 12.10 13.80 7.6±0.7 3.4±0.6 2.0 2.24±0.44

G82.2+5.3 82.2 5.3 6.87 8.30 4.7±0.5 2.8±0.1 2.0 1.68±0.19

G85.9-0.6 85.9 -0.6 7.59 9.45 6.9±0.3 0.7±0.1 5.0 9.86±1.47

G89.0+4.7 northwest 89.0 4.7 6.83 7.64 3.0±0.4 1.6±0.3 1.7 1.88±0.43

G109.1-1.0 109.1 -1.0 6.39 7.50 7.1±0.7 9.0±1.4 3.1±0.6 3.2 2.29±0.50 2.90±0.72

G111.7-2.1 north 111.7 -2.1 3.94 6.70 13.6±0.3 11.7±1.6 4.8±0.7 3.4 2.83±0.42 2.44±0.49

G116.9+0.2 116.9 0.2 6.32 6.77 5.9±0.8 9.2±0.7 2.6±0.1 3.0 2.27±0.32 3.54±0.30

G119.5+10.2 119.5 10.2 2.06 2.17 2.1±0.4 3.8±1.1 1.3±0.2 1.4 1.62±0.40 2.92±0.96

G120.1+1.4 120.1 1.4 8.24 9.22 5.8±0.4 8.0±1.0 1.8±0.1 3.0 3.22±0.29 4.44±0.61

G130.7+3.1 130.7 3.1 5.57 5.72 4.1±0.1 5.4±0.1 2.1±0.2 2.0 1.95±0.19 2.57±0.25

G132.7+1.3 132.7 1.3 7.98 9.00 8.3±0.3 2.2±0.1 2.5 3.91±0.22

G166.0+4.3 166.0 4.3 3.86 5.67 2.0±0.2 1.7±0.2 4.5 1.18±0.18

G180.0-1.7 180.0 -1.7 4.93 6.30 3.1±0.2 0.7±0.2 1.3 4.43±1.30

G184.6-5.8 184.6 -5.8 3.26 3.79 3.2±0.1 4.1±0.1 1.5±0.1 2.0 2.13±0.16 2.73±0.19

G189.1+3.0 189.1 3.0 4.85 6.09 5.2±0.5 6.0±0.5 2.8±0.6 1.5 1.86±0.44 2.14±0.49

G260.4-3.4 260.4 -3.4 7.45 9.46 3.1±0.1 3.5±1.5 1.5±0.2 2.2 2.07±0.28 2.33±1.05

G263.9-3.3 263.9 -3.3 7.38 9.10 0.19±0.02 0.3±0.1 0.31±0.05 0.29 0.61±0.12 0.97±0.36

G292.0+1.8 292 1.8 5.52 6.21 3.8±0.3 5.0±0.4 2.2±0.2 6.0 1.73±0.21 2.27±0.28

G296.1-0.5 north 296.1 -0.5 13.20 16.20 0.6±0.2 1.9±0.3 2.0 0.32±0.12

G296.5+10.0 296.5 10.0 1.13 1.27 1.2±0.1 0.5±0.1 2.1 2.40±0.52

G315.4-2.3southwest 315.4 -2.3 7.19 9.06 4.0±0.2 1.7±0.2 2.3 2.35±0.30

G320.4-1.2 320.4 -1.2 14.80 14.70 10.9±0.1 11.5±0.3 3.6±0.4 5.2 3.03±0.34 3.19±0.36

G327.6+14.6 327.6 14.7 0.75 0.68 0.68±0.10 1.6±0.1 0.32±0.10 2.2 2.13±0.73 5.00±1.59

G332.4-0.4 332.4 -0.4 20.70 22.80 7.0±0.7 8.7±3.5 4.7±0.7 3.1 1.49±0.27 1.85±0.79

G332.5-5.6 332.5 -5.6 2.47 2.84 1.4±0.4 0.8±0.3 3.0 1.75±0.83

Planetary nebulae

G37.7-34.5 37.7 -34.5 0.50 0.47 0.81±0.85 0.19±0.02 1.1 4.26±4.50

G64.7+5.0 64.7 5.0 3.03 3.10 2.5±0.1 1.0±0.1 1.2 2.50±0.27

G84.9-3.4 84.9 -3.4 4.68 5.58 4.1±0.2 2.9±0.3 0.7 1.41±0.16

G94.0+27.4 94.0 27.4 0.34 0.41 0.12±0.02 0.09±0.03 1.0 1.33±0.50

G197.8+17.3 197.8 17.3 0.54 0.57 0.8±0.6 0.5±0.1 1.2 1.60±1.24

G331.7-1.0 331.7 -1.0 16.70 17.60 5.0±0.8 4.1±0.4 1.1 1.22±0.23

X-ray Binaries

2E 0236.6+6101 135.7 1.1 7.57 9.02 5.9±0.2 2.8±0.3 2.4 2.11±0.24

2E 1118.7-6138 292.5 -0.9 12.00 14.10 8.1±0.1 13.2±0.3 2.9±0.3 5.0 2.79±0.29 4.55±0.48

2S 0620-003 209.6 -6.5 3.66 4.10 2.4±1.1 1.2±0.1 1.2 2.00±0.93

2S 0921-63 281.8 -0.9 1.70 2.13 2.3±0.1 1.02±0.03 8.5 2.25±0.12

4U 0352+30 163.1 -17.1 7.09 9.39 2.5±1.1 1.2±0.1 1.0 2.08±0.19

4U 0538+26 181.4 -2.6 4.50 5.90 3.9±0.1 7.0±0.3 2.4±0.2 2.5 1.63±0.08 2.92±0.17

4U 0614+09 200.9 -3.4 4.48 5.47 3.5±0.1 3.3±0.1 1.2±0.1 2.92±0.26 2.75±0.24

4U 0919-54 275.9 -3.8 6.16 7.30 3.1±0.1 3.0±0.1 1.0±0.1 4.8 3.10±0.33 3.00±0.32

4U 1036-56 285.4 1.4 9.02 10.10 2.8±0.9 2.3±0.8 5.0 1.22±0.58

4U 1118-60 292.1 0.3 9.95 12.00 11.0±0.1 4.0±0.5 9.0 2.75±0.34

4U 1258-61 304.1 1.2 10.9 10.6 10.0±0.2 5.5±0.7 2.4 1.82±0.23

4U 1254-69 303.5 -6.4 2.16 2.91 2.6±0.1 2.2±0.1 1.2±0.1 13 2.17±0.55 1.83±0.47

4U 1456-32 332.2 23.9 0.91 0.84 0.7±0.1 0.9±0.1 0.31±0.15 1.2 2.26±1.14 2.90±1.44

4U 1516-56 322.1 0.0 15.90 19.80 18.4±1.2 12.2±1.6 5.5 1.51±0.22

4U 1538-52 327.4 2.2 9.14 9.70 16.2±0.5 20.6±0.5 7.1±0.7 5.5 2.28±0.24 2.90±0.29

4U 1543-47 330.9 5.4 3.51 4.00 3.8±0.2 1.6±0.2 7.5 2.38±0.32

4U 1543-62 321.8 -6.3 2.43 2.96 3.2±0.1 3.4±0.1 1.0±0.3 7.0 3.20±0.97 3.40±1.02

4U 1608-52 330.9 -0.9 18.10 19.20 10.7±0.2 10.8±2.2 6.2±1.7 3.9 1.73±0.47 1.74±0.60

4U 1636-53 332.9 -4.8 2.76 3.58 2.9±0.1 3.7±0.1 2.5±0.3 6.0 1.16±0.14 1.48±0.18

4U 1658-48 338.9 -4.3 3.74 5.26 5.0±0.1 8.1±0.2 3.7±0.4 1.35±0.15 2.19±0.24

4U 1704-30 353.8 7.3 1.82 1.82 2.3±0.2 0.9±0.1 10.0 2.56±0.36

4U 1724-307 356.3 2.3 5.70 6.57 6.6±0.5 4.9±0.6 7.4 1.35±0.19

c© 0000 RAS, MNRAS 000, 000–000


4U 1728-16 8.5 9.0 1.96 2.08 2.4±0.1 0.8±0.5 4.2 3.00±1.88

4U 1728-34 354.3 0.2 12.40 16.00 29.3±0.5 10.8±1.8 5.2 2.71±0.45

4U 1730-335 354.8 -0.2 12.20 16.70 17.4±0.3 10.1±1.9 8.0 1.72±0.19

4U 1746-37 353.5 -5.0 2.63 2.95 4.0±0.3 1.5±0.2 11.0 2.67±0.41

4U 1820-30 2.8 -7.9 1.29 1.50 2.7±0.1 2.6±0.3 0.9±0.1 7.6 3.00±0.35 2.89±0.46

4U 1822-37 356.9 -11.3 1.04 1.24 1.2±0.2 0.5±0.3 2.5 2.40±1.49

4U 1837+04 36.1 4.8 4.14 4.71 5.2±0.2 1.3±0.1 8.4 4.00±0.34

4U 1850-08 25.4 -4.3 2.36 2.78 3.8±0.1 4.8±0.2 1.2±0.2 8.2 3.17±0.53 4.00±0.69

4U 1908+00 35.7 -4.1 2.84 3.37 4.4±0.1 4.1±0.1 1.2±0.1 5.0 3.67±0.32 3.42±0.30

4U 1915-05 31.4 -8.5 2.31 2.60 3.9±0.1 5.6±40.7 2.2±0.9 8.9 1.77±0.73 2.55±1.09

4U 1956+35 71.3 3.1 7.21 7.21 4.4±0.1 5.4±0.1 3.3±0.1 2.14 1.33±0.05 1.64±0.06

4U 1957+11 51.3 -9.3 1.21 1.27 1.6±0.1 1.1±0.1 0.9±0.1 1.78±0.23 1.22±0.18

4U 2129+47 91.6 -3.0 3.30 3.82 2.4±0.7 3.6±0.1 1.2±0.2 7.0 2.00±0.67 3.00±0.51

4U 2135+57 99 3.3 8.18 9.19 7.8±0.3 4.0±0.3 3.8 1.95±0.16

4U 2142+38 87.3 -11.3 1.87 2.17 1.9±0.1 2.9±0.1 1.1±0.1 7.2 1.73±0.18 2.64±0.26

EXO 0331+530 146.1 -2.2 6.91 8.54 9.0±1.0 5.8±0.2 1.55±0.18

EXO 0748-676 280.0 -19.8 1.01 1.14 0.8±0.1 0.6±0.1 0.2±0.1 6.8 4.00±2.06 3.00±1.58

EXO 1745-248 3.8 1.7 6.33 5.81 13.2±3.5 7.2±1.0 8.7 1.83±0.55

EXO 2030+375 77.2 -1.2 8.80 10.30 20.5±0.1 11.7±0.7 7.1 1.75±0.11

GRO J1719-24 0.1 7.0 2.56 2.84 4.0±0.4 2.8±0.6 1.2 1.43±0.34

GRO J0422+32 165.9 -11.9 1.55 1.66 1.6±0.9 0.8±0.1 2.00±1.15

GRO J1008-57 283 -1.8 13.50 15.10 12.7±0.1 6.1±0.2 2.0 2.08±0.07

GRO J1655-40 345 2.5 5.33 6.86 7.4±0.3 7.4±0.2 3.7±0.3 2.00±0.18 2.00±0.17

GRO J1944+26 63.2 1.4 9.43 10.20 13.0±0.1 16.3±0.3 6.5±0.3 9.5 2.00±0.09 2.51±0.12

GRO J2058+42 83.5 -2.7 6.22 7.1 9.5±2.3 4.3±0.4 9.0 2.21±0.57

GRS 1009-45 275.9 9.3 1.05 1.16 0.4±0.1 0.6±0.2 4.5 0.67±0.28

GS 0834-43 262.0 -1.5 10.10 11.80 21.2±2.3 12.3±3.7 4.0 1.72±0.29

Gu Mus 295.3 -7.1 1.91 2.50 2.3±0.1 0.9±0.1 5.9 2.56±0.30

IGR J17544-2619 3.2 -0.3 11.80 14.40 11.2±1.2 6.3±0.4 3.7 1.78±0.22

IGR J22534+6243 110.0 2.9 8.96 9.98 18.1±1.9 5.7±0.4 3.18±0.40

KS 1947+300 66.1 2.1 8.90 10.40 5.1±0.2 3.4±0.4 10.0 1.50±0.19

LS 2883 304.2 -1.0 14.50 16.20 4.0±0.3 3.3±0.4 2.7 1.21±0.17

LS 5039 16.9 -1.3 8.87 8.71 6.4±0.1 3.7±0.3 2.9 1.73±0.14

ALS 19596 33.0 1.7 10.10 10.10 23.0±1.3 8.1±0.9 2.84±0.35

MXB 1746-20 7.7 3.8 3.30 2.92 6.7±0.3 5.7±0.1 3.2±0.3 8.5 2.13±0.22 1.78±0.17

SAX J1748.9-2021 7.7 3.8 3.30 2.92 7.0±1.0 3.2±0.4 8.5 2.19±0.42

RX J0146.9+6121 129.5 -0.8 7.58 8.49 4.9±0.5 2.9±0.1 2.3 1.69±0.18

RX J0440.9+4431 159.8 -1.3 5.45 6.64 3.4±0.1 2.2±0.2 3.3 1.55±0.15

RX J1829.4-2347 9.3 -6.1 1.71 1.88 2.6±1.7 1.2±0.4 6.0 2.17±0.73

RX J1832-33 1.5 -11.4 0.92 1.23 0.55±0.22 0.58±0.06 0.3±0.1 9.3 1.83±0.95 1.93±0.67

Swift J1753.5-0127 24.9 12.2 1.74 1.69 2.2±0.1 2.7±0.1 1.3±0.1 6.0 1.69±0.15 2.08±0.18

V404 Cyg 73.1 -2.1 6.66 8.10 8.1±0.3 8.6±0.6 3.6±0.5 3.5 2.25±0.32 2.39±0.37

XTE J1550-564 325.9 -1.8 10.10 8.97 6.6±0.1 6.6±0.6 2.9±0.2 5.3 2.28±0.16 2.28±0.16

XTE J1720-318 354.6 3.1 4.77 5.25 15.1±0.2 6.6±0.6 2.29±0.21

XTE J1808-369 355.4 -8.1 1.13 1.34 1.4±0.1 1.4±0.1 0.8±0.3 3.0 1.75±0.67 1.75±0.67

XTE J1819-254 6.8 -4.8 1.81 2.34 2.4±0.1 0.8±0.1 9.5 3.00±0.63

XTE J1859+226 54.0 8.6 2.16 2.21 3.6±0.4 1.8±0.4 7.6 2.00±0.50

XTE J2103+457 87.1 -0.7 6.64 7.78 7.6±0.1 4.2±0.3 6.5 1.81±0.13

XTE J2123-058 46.5 -36.2 0.54 0.57 0.8±2.3 0.37±0.15 8.5 2.16±6.28

a Gl: Galactic longitude; Gb: Galactic latitude.b KBH: HI column density at the source (column one) direction based on Kalberla, Burton, Hartmann et al. (2005).c DL: HI column density at the source (column one) direction based on Dickey & Lockman (1990).d NH(AG89): NH derived with the solar abundances of AG89.e NH(W00): NH derived with subsolar abundances (see Wilms et al. 2000 and references therein).f D: Distance.

et al. 2001), [SII](∼10320 A)/[SII](∼4068 A) (e.g., SNR G111.7-

2.1, Hurford & Fesen 1996) and [FII](∼16435 A)/[FII](∼12567 A)

(e.g., SNR G332.4-0.4, Oliva et al. 1989). However, caution needs

to be taken when applying the Balmer decrement method to very

young SNRs, because their intrinsic Hα/Hβ ratios may be different

from the canonical ratios used for nebulae under the “Case B” con-

dition, in which the nebula is always optical thin except for Lyman-

line photons (e.g., Tycho, Ghavamian et al. 2000). Other methods to

estimate the extinction of SNRs include the use of existing extinc-

tion maps (e.g., SNR G119.5+10.2, Mavromatakis et al. 2000), and

the association of the targets with extinction-known objects (e.g.,

SNR G260.4-3.4, Gorenstein 1975).

For XBs, four methods are usually adopted to estimate their

extinction: (1) the relation between the colour excess or redden-

ing E(B − V ) and the widths of the diffuse interstellar bands or

Na doublet lines (e.g., 4U0614+091, Nelemans et al. 2004); (2)

the colour-magnitude diagram for those in star clusters (e.g., 4U

1820-30 in NGC 6624, Valenti et al. 2004); (3) a comparison of the

observed colour with the intrinsic colour derived from the stellar

spectral type (e.g., 4U1908+005, Thorstensen et al. 1978); and (4)

fitting the observed, reddened UV spectrum of a source (e.g., 2S

0620-003) to determine the excess extinction at 2175 A (∆Abump)

and assuming a constant ∆Abump/AV ratio.

For each source in our sample we estimate its AV or E(B −V )1 following these rules: (1) if a source has more than one mea-

sure of AV or E(B−V ), an error-weighted value will be adopted;

(2) if a source has more than two measures of AV or E(B − V )and one measure differs from the others by >

∼30% or 3σ, this mea-

sure will be rejected; (3) if there is no error estimation for the AV

or E(B − V ) measure, we will assign an uncertainty of ∼ 13%

which is obtained by averaging over those sources for which the

uncertainties on AV or E(B − V ) are known (see Column 13 in

Table 7).

1 From E(B − V ) to AV we take AV = 3.1×E(B − V ).

c© 0000 RAS, MNRAS 000, 000–000


2.2 NH

The X-ray absorption factor is usually expressed as

M(E) = exp {NHσeff(E)} , (1)

where σeff(E) represents the effective absorption cross section tak-

ing into account the composition of interstellar medium. Appar-

ently, the interstellar X-ray absorption is affected by both NH and

σeff(E). To be self-consistent, it is necessary to seek NH from the

literature using similar heavy element abundances (e.g., O, Ne, Fe,

Mg, Si). As mentioned earlier, different abundance selection (e.g.,

AG89 vs W00) may lead to a difference of > 30% for the derived

NH (Foight et al. 2016). In this study, we will consider the hydro-

gen column densities derived either from solar abundances (AG89)

or subsolar abundances (W00). If there is no mention in the litera-

ture about the abundance setting for the software XSPEC, we will

assume that the AG89 abundance was used because it is the default

setting at least back to XSPEC version 8.

We take the following approach for adopting NH: (1) if a

source has more than one measure of NH, an error-weighted value

will be adopted; (2) if there are more than one measure of NH and

these measures all differ from each other by >∼30% or 3σ, the one

derived from the data with a higher quality at energy E<∼2 keV

will be selected because this energy range is more sensitive to X-

ray absorption; this means that those derived from Chandra, XMM-

Newton and Suzaku have a higher priority than those from ROSAT,

RXTE and other telescopes; (3) if there are two or three measures

of NH for the same source and they are all derived from Chan-

dra, XMM-Newton or Suzaku, we will take the error-weighted value

without considering their differences; (4) if there is no error estima-

tion for a NH measure, we will adopt an uncertainty of ∼ 9% for

AG89 or ∼ 8% for W00 which are obtained by averaging over those

sources for which the uncertainties on NH are known (see Columns

5, 9 in Table 7); (5) for SNRs which are extended sources of which

AV and NH are known for a range of regions, we make sure that

the adopted AV and NH are from the same region, if possible. Note

that for XBs, some of them are dipper sources and have intrinsic lo-

cal absorptions. We try to avoid these local absorptions by taking

the following procedure: (1) for low mass X-ray binaries (LMXBs),

we examine the catalog of Liu et al. (2006) to see whether they are

dipper sources or not; if yes, we will exclude them from our sample

except those for which NH was derived from the spectral fitting and

no dipper shows up in the spectra or the interstellar NH and local

NH are considered separately during the fitting (e.g., using the par-

tial covering model); (2) for high mass X-ray binaries (HMXBs),

we examine them one by one through the literature to see whether

the interstellar NH and local NH are considered separately and

whether the interstellar NH varies during different phases and states

to avoid the influence of local absorptions. With these selection

procedures, seven XBs (4U 0115+63, HD 259440, IGR J11215-

5952, IGR J18450-0435, IGR J2000.6+3210, IGR J21343+4738,

and LS 992) are exclused due to the possible presence of local ab-

sorptions. For three XBs (GRO J2058+42, GS 0834-43, and ALS

19596), we find no evidence for or against the presence of con-

siderable amounts of local absorption. As their NH/AV ratios are

not much larger than those previously published, we decide to keep

them in our sample. We also note that, although SNRs are treated

as having little or no intrinsic absorption, some of them may actu-

ally have considerable local dust emission (e.g., Cas A [Dunne et

al. 2009]).

In Table 4, we present the finally adopted data. For each source

we also list the NH value determined from the HI 21 cm observa-

Table 5. A summary of the final NH-AV fitting.

Sample Coefficient Zero Offset

(1021 Hcm−2 mag−1) (1021 Hcm−2)

AG89

Whole 2.08±0.02 0 (fixed)

Whole 2.12±0.03 −0.11± 0.04SNRs 2.06±0.05 0 (fixed)

XBs 2.09±0.03 0 (fixed)

W00

Whole 2.47±0.04 0 (fixed)

Whole 2.59±0.07 −0.23± 0.10SNRs 2.51±0.10 0 (fixed)

XBs 2.46±0.05 0 (fixed)

tion toward their lines of sight. One can find all the compiled NH

and AV for each sources in Appendix.

3 RESULTS AND DISCUSSION

As the AG89 sample (i.e., those sources for which NH was derived

from solar abundances) is nearly three times larger than the W00

sample (i.e., those sources for which NH was derived from subso-

lar abundances), the former sample allows us not only to determine

NH/AV but also to investigate the variation of NH/AV and the

distribution of interstellar gas in our Galaxy. Therefore, in this sec-

tion, we will first focus on the AG89 sample in §3.1–§3.4. We will

consider the W00 sample in §3.5. The possible caveats of our study

will be discussed in §3.6.

3.1 The mean NH–AV relation

Figure 1 plots NH against AV for all calibration points with

abundance of AG89. We fit a linear relationship to the data using

the MPFITXY routine (Williams, Bureau & Cappellari 2010) of

the MPFIT package (Markwardt 2009):

NH(H cm−2) = (2.19 ± 0.02) × 1021 AV (mag), (2)

where the errors represent a 1σ statistical uncertainty. The

black solid lines in Figure 1 shows the best-fit result. To assess

the influence of possible intrinsic absorption in our sample,

we add a zero point offset to equation (2) and we then obtain a

slightly steeper relation as shown as the red dashed line in Figure 1.

NH(H cm−2) = (2.25± 0.03) × 1021 AV (mag)− (0.11 ± 0.04) × 1021(H cm−2) .

(3)

Since the zero point offset in equation (4) is much smaller than

the column density of most of our calibration points, it indicates

that our sample is not affected significantly by the possible intrinsic

absorption.

We also explore whether different kinds of calibrator affect the

derived NH-AV relation. To this end, we fit the data of SNRs and

XBs separately and obtain

NH(H cm−2) = (2.51±0.05)×1021 AV (mag) (4)

for SNRs, and

NH(H cm−2) = (2.09±0.03)×1021 AV (mag) (5)

for XBs. Surprisingly, they are inconsistent with each other and

c© 0000 RAS, MNRAS 000, 000–000


Figure 4. The distribution of NH/AV of the AG89 sample projected on the Galactic longitude-latitude plane (upper panel) and the Galactic disk (lower panel).

In the lower panel, the Galactic centre is marked by a plus sign. The two dot circles indicate 2 kpc and 10 kpc away from the Galactic centre. The solid lines

present the spiral arms from Hou & Han (2014) and Green et al. (2011).

c© 0000 RAS, MNRAS 000, 000–000


Figure 5. Variations of NH/AV with Rg , z, nH and NH.

Figure 6. NH–AV relations for the Galactic 1st/4th quadrants (left) and for the Galactic 2nd/3rd quadrants (right).

c© 0000 RAS, MNRAS 000, 000–000


the mean 〈NH/AV 〉 ratio derived from SNRs is larger by ∼ 20%

than that from XBs. To find what causes this difference, we

show the NH–AV relations in Figure 2 for SNRs and XBs, sep-

arately. Apparently, those three SNRs (G54.1+0.3, G85.9-0.6, and

G132.7+3.1) which lie well above the best-fit line (see Figure 2,

right panel) lead to an overestimation of NH/AV . Indeed, for

G54.1+0.3, Kim et al. (2013) estimated AV ∼ 7.3±0.1mag based

on the extinction toward several stars associated with the SNR,

while the extinction toward the majority of the stars is in the range

from ∼ 6.9 to ∼ 7.5mag and for one star AV is as high as ∼ 7.8

mag (see their Table 3). Also, Koo et al. (2008) derived an extinc-

tion of 8.0±0.7mag. Therefore, an uncertainty of 0.1mag may

have been underestimated. For G85.9-0.6, Gok et al. (2009) ob-

tained AV =0.47±0.05 mag, 1.18±0.12mag, and 1.40±0.14magat three positions which are not fully overlapped with the X-ray

bright region. Therefore, the error-weighted AV and the associated

uncertainty that we used may be not a good estimation for the X-

ray bright region. For G132.7+3.1, the optical spectroscopy was

taken at the western filaments (Fesen et al. 1995). However, only

its central and northern regions are bright at X-ray. While the size

of G132.7+3.1 is about 80′, the space separation between the op-

tical and the X-ray observations could be larger than 0.5o. Taking

these concerns into consideration, we decide to exclude these three

SNRs. We then obtain

NH(H cm−2) = (2.06±0.05)×1021 AV (mag) (6)

for SNRs, and

NH(H cm−2) = (2.08±0.02)×1021 AV (mag) (7)

for the whole AG89 sample. We present the best-fit lines in Fig-

ure 3. Table 5 summarizes the above results. Our results on the

NH/AV ratio with AG89 abundance are compatible with those pre-

viously obtained from X-ray data with the same abundance.

3.2 A constant NH/AV ratio across the Galaxy?

Watson (2011) compared the HI column density with the extinction

with data from the LAB HI survey (Kalberla et al. 2005) and the ex-

tinction map of Schlegel et al. (1998). He found that the NH/AV

ratio seems to correlate with NH toward the Galactic centre. To

investigate the variation of NH/AV within our Galaxy, we plot

the NH/AV ratio distribution projected on the Galactic longitude-

latitude (l-b) plane and the Galactic disk in Figure 4. There is no

clear indication for a large-scale variation of NH/AV within the

Galaxy. In Figure 5 we show the NH/AV ratio against (i) Rg , the

distance away from the Galactic centre, (ii) z, the distance away

from the Galactic plane, (iii) nH, the average line-of-sight number

density of hydrogen, and (iv) NH. It is apparent that the NH/AV

ratio does not show any systematic variation with Rg , z, nH or NH.

We also investigate the NH–AV relation for the inner Galaxy (the

Galactic first/fourth quadrants) and outer Galaxy (the second/third

quadrants) separately. We derive

NH(H cm−2) = (2.04±0.05)×1021 AV (mag) (8)

for the first/fouth quadrants, and

NH(H cm−2) = (2.09±0.03)×1021 AV (mag) (9)

for the second/third quadrants (see Figure 6). They are consistent

within each other. In summary, we can safely conclude that the X-

ray derived NH/AV ratio is roughly a constant across the entire

Galaxy.

Figure 7. The projection distribution of the AG89 sample on the Galactic

plane. The Galactic center is marked by a plus sign. The dot circles and

solid lines are the same as the circles and lines in Figure 4.

We should note that it has been well known that the abun-

dances of heavy elements in our Galaxy increase toward the Galac-

tic centre (e.g., see Rolleston et al. 2000). Since AG89 assumed a

uniform abundance for the ISM, at the Galactic centre NH could

have been overestimated while at the outer Galaxy NH could have

been underestimated. However, it is also well recognized that the

dust grains in the Galactic centre could be much larger than that

of the diffuse ISM (e.g., see Gao et al. 2010, Shao et al. 2017)

and the visual extinction on a per unit mass basis could be smaller.

Therefore, even the Galactic centre is rich in dust-forming heavy

elements and more heavy elements have been tied up in dust grains

one does not necessarily expect a lower NH/AV .

3.3 The Gas-to-dust mass ratio

From the hydrogen-to-extinction ratio we can determine

Mgas/Mdust, the gas-to-dust mass ratio:

Mgas

Mdust= 1.086 × 1.4µH κext(V )

(

NH

AV

)

, (10)

where µH is the atomic mass of hydrogen and κext(V ) is the V -

band mass extinction coefficient (i.e., extinction cross section per

dust mass). The coefficient, 1.4, arises from the contribution of he-

lium. Taking κext(V ) from Li & Draine (2001) and Draine & Li

(2007), we obtain Mgas/Mdust ≈ 140.

3.4 The gas distribution of the Galaxy

When both NH and D (the distance of the source to Earth) are

known, we are able to study the distribution of the gas in the

Galaxy. To this end, we enlarge the sample in Table 1 by including

74 SNRs with known NH (AG89 abundance) and D from the litera-

ture (see Table 6). The final sample includes 175 calibrators. Figure

7 shows the projection distribution of the sample on the Galactic

plane.

The distribution of hydrogen away from the Galactic plane

is usually described by one to three Gaussian and/or exponen-

tial functions, (e.g., see Bohlin, Savage & Drake 1978, Diplas

& Savage 1994, Dickey & Lockman 1990, Kalberla & Kerp

2009). A commonly-used HI distribution between the Galactic

c© 0000 RAS, MNRAS 000, 000–000


Table 6. Hydrogen column density NH and distance D compiled from the literature.

Objects Gl Gb NH D References

Name (degree) (degree) (1021 Hcm−2) (kpc)

G1.0-0.1 1.0 -0.1 75 8.5 1,2

G1.9+0.3 1.9 0.3 58 8.5 3,4

G5.4-1.2 5.4 -1.2 35 5.2 1,5

G8.7-0.1 8.7 -0.1 12 4.5 6

G12.0-0.1 12.0 -0.1 49 10.9 7,8

G12.8-0.0 12.8 0.0 100 4.8 7,9

G15.9+0.2 15.9 0.2 39 8.5 7,10

G16.7+0.1 16.7 0.1 47.4 14.0 11,12

G18.9-1.1 18.9 -1.1 8.3 2.0 13,14

G21.5-0.9 21.5 -0.9 22.4 4.8 15,16

G21.8-0.6 21.8 -0.6 28 5.5 15,17

G27.4+0.0 27.4 0.0 26 8.7 18,19

G28.6-0.1 28.6 -0.1 37 7.0 20,21

G28.8+1.5 28.8 1.5 20 4.0 22

G29.7-0.3 29.7 -0.3 29 10.6 23-30

G31.9+0.0 31.9 0.0 28 7.9 31-36

G32.4+0.1 32.4 0.1 52 17.0 37

G32.8-0.1 32.8 -0.1 8.1 4.8 38,39

G33.6+0.1 33.6 0.1 15 7.5 40-44

G34.7-0.4 34.7 -0.4 13 2.7 43,45-48

G41.1-0.3 41.1 -0.3 32 10.3 49,50

G43.3-0.2 43.3 -0.2 50 10.0 51-55

G49.2-0.7 49.2 -0.7 17 5.2 56-61

G65.7+1.2 65.7 1.2 2 0.9 62-64

G76.9+1.0 76.9 1.0 16 10.0 65,66

G85.4+0.7 85.4 0.7 8.3 3.8 67,68

G148.1+0.8 148.1 0.8 4.9 1.0 69-71

G156.2+5.7 156.2 5.7 3.2 2.0 72-75

G195.1+4.3 195.1 4.3 0.11 0.25 71,76

G201.1+8.3 201.1 8.3 0.25 0.28 71,77-79

G266.2-1.2 266.2 -1.2 3.5 0.75 80-84

G272.2-3.2 272.2 -3.2 10 10.0 85-89

G284.0-1.8 294.0 -1.8 5 3.0 90,91

G284.2-0.4 284.2 -0.4 15 9.0 92,93

G284.3-1.8 284.3 -1.8 6.6 2.9 94,95

G287.4+0.6 287.4 0.6 9 3.0 96,97

G290.1-0.8 290.1 -0.8 7.3 7.0 98,99

G291.0-0.1 291.0 -0.1 6.7 3.5 100,101

G292.2-0.5 292.2 -0.5 14.6 8.4 102-104

G296.7-0.9 296.7 -0.9 12.4 9.8 105

G296.8-0.3 296.8 -0.3 6.4 9.6 106,107

G299.2-2.9 299.2 -2.9 3.5 5.0 108,109

G304.1-0.2 304.1 -0.2 27 7.0 93,110

G304.6+0.1 304.6 0.1 38.7 10.0 89,111,112

G306.3-0.9 306.3 -0.9 19.5 8.0 113

G308.3-1.4 308.3 -1.4 10 9.8 114,115

G309.2-0.6 309.2 -0.6 6.5 4.0 116

G309.8-2.6 309.8 -2.6 3.9 2.5 117,118

G310.6-1.6 310.6 -1.6 20.9 10.0 93,119

G311.5-0.3 311.5 -0.3 27.7 13.8 89,112,120

G313.3+0.1 313.3 0.1 18.5 3.0 93,121

G313.6+0.3 313.6 0.3 35.4 5.6 121-123

G326.3-1.8 326.3 -1.8 3.5 5.1 124,125

G327.1-1.1 327.1 -1.1 18.8 9.0 126-128

G327.4+0.4 327.4 0.4 24 4.3 129,130

G330.2+1.0 330.2 1.0 26 7.8 89,130,131

G337.2+0.1 337.2 0.1 83 11.0 132-134

G337.2-0.7 337.2 -0.7 32.3 9.0 89,135

G337.8-0.1 337.8 -0.1 67.5 12.3 136,137

G338.3-0.0 338.3 0.0 150 12.8 138,139

G343.7-2.3 343.7 -2.3 5 2.6 71,140,141

G344.7-0.1 344.7 -0.1 45 6.3 142

G346.6-0.2 346.6 -0.2 22.7 11.0 120,137,143-

146

c© 0000 RAS, MNRAS 000, 000–000


G347.3-0.5 347.3 -0.5 7.8 1.0 146-149

G348.5+0.1 348.5 0.1 35 7.9 120,150,151

G348.7+0.3 348.7 0.3 38 13.2 151-153

G349.7+0.2 349.7 0.2 64 11.5 154,155

G350.1-0.3 350.1 -0.3 33 9.0 154

G352.7-0.1 352.7 -0.1 32 7.6 156,157

G353.6-0.7 353.6 -0.7 18.2 3.2 158-160

G355.6-0.0 355.6 0.0 55 13.0 145,161

G357.7-0.1 357.7 -0.1 100 11.8 31,162

G359.0-0.9 359.0 -0.9 16 3.7 89,163,164

G359.1-0.5 359.0 -0.5 25 7.6 163-165

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& Reynolds (1994); (53) Hwang et al. (2000); (54) Brogan & Troland (2001); (55) Keohane et al. (2007); (56) Green et al. (1997); (57) Koo et al. (2002); (58)

Tian & Leahy (2013); (59) Koo et al. (2005); (60) Hanabata et al. (2013); (61) Sasaki et al. (2014); (62) Kothes et al. (2004); (63) Arzoumanian et al. (2004);

(64) Arzoumanian et al. (2008); (65) Marthi et al. (2011); (66) Arzoumanian et al. (2011); (67) Kothes et al. (2001); (68) Jackson et al. (2008); (69) McGowan

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et al. (2014b); (151) Tian & Leahy (2012); (152) Nakamura et al. (2009); (153) Aharonian et al. (2008a); (154) Yasumi et al. (2014); (155) Tian & Leahy

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radius 4 kpc and 8 kpc is given by Dickey & Lockman (1990):

nH(z) = nH,1(0) exp(

−z2/2h21

)

+ nH,2(0) exp(

−z2/2h22

)

+nH,3(0) exp

(

−z2/2h23

)

, where the midplane densities are

nH,1(0) ≈ 0.395 cm−3, nH,2(0) ≈ 0.107 cm−3, and nH,3(0) ≈0.064 cm−3, and the scale heights are h1 ≈ 90 pc, h2 ≈225 pc, and h3 ≈ 403 pc. A more recent study for the Galac-

tic radius between 5 kpc and 35 kpc suggested that the HI scale

hight varies with the distance away from the Galactic cen-

tre: h(Rg) ≈ 150 exp {(Rg −R0)/9800} pc, where R0 is

the Galactic radius of the Sun (Kalberla & Kerp 2009). Fig-

ure 8 shows NH sin(|b|) versus z. To avoid the influence of the

Galactic bulge and wrap, we only make use of the data within

2 kpc<∼Rg<∼10 kpc. We use a one-component Gaussian model,

nH(z) = nH(0) exp(

−z2/2h2)

, to fit the data. The best-fit is

given by nH(0) = 1.11± 0.15 cm−3 and h = 75.5± 12.4 pc (see

Figure 8). Compared with that of Dickey & Lockman (1990), our

results suggest a larger midplane density and a sightly smaller scale

height. This might be caused by the fact that the NH column den-

sities adopted here are based on X-ray absorption data and contain

the contribution of molecular and ionized ISM, while the 21 cm

emission only presents atomic hydrogen. We are not able to esti-

mate the scale heights of any other extended components, because

the majority of our sample has z < 1000 pc.

The neutral hydrogen surface density is nearly constant within

the Galactic radius of ∼ 12.5 kpc (Kalberla & Kerp 2009). A mean

surface density of ∼ 5M⊙ pc−2 was given by Dickey & Lock-

man (1990). More recently, Kalberla & Dedes (2008) argued for a

surface density two times larger than that of Dickey & Lockman

(1990). Based on the best-fit nH(0) and h values, here we estimate

a mean gas surface density of the Galaxy in the Galactic radius

range of ∼ 2 kpc to ∼ 10 kpc to be∑

gas ∼ 1.4nH(0) × 2.4 h ≈

7.0M⊙ pc−2.

Due to the limited amount of data, we are not able to accu-

rately determine the radial distribution of the midplane density, but

only give an rough estimation of nH ≈ 4.33 ± 12.87 cm−3 for

2.0 kpc < Rg < 4.5 kpc, nH ≈ 1.15 ± 5.35 cm−3 for 4.5 kpc <Rg < 6.5 kpc, nH ≈ 1.18 ± 0.77 cm−3 for 6.5 kpc < Rg <8.5 kpc, and nH ≈ 1.13±0.81 cm−3 for 8.5 kpc < Rg < 10 kpc.

All regions are limited to |z| < 75 pc.

3.5 NH/AV from the W00 Sample

Since the X-ray-derived NH is sensitive to the abundances of heavy

elements (relative to H), one could expect a higher NH/AV for the

W00 sources of which NH is determined with subsolar abundances

c© 0000 RAS, MNRAS 000, 000–000


Figure 8. Logarithm of NH sin |b| vs. the logarithm of z (distance above

the Galactic plane). The solid line is the best fit to the data with a midplane

density of nH(0) = 1.11 ± 0.15 cm−3 and a scale hight of h = 75.5 ±12.4 pc.

compared to that of the AG89 sources of which NH is determined

with solar abundances. By fitting the whole W00 sample which in-

clude both SNRs and XBs, we indeed obtain a larger 〈NH/AV 〉:

NH(H cm−2) = (2.87 ± 0.04) × 1021 AV (mag) . (11)

We also obtain

NH(H cm−2) = (3.30±0.06)×1021 AV (mag) (12)

for the SNRs of the W00 sample, and

NH(H cm−2) = (2.46±0.05)×1021 AV (mag) (13)

for the XBs of the W00 sample. Again, the mean 〈NH/AV 〉 ratio

from SNRs differs from that of XBs by ∼ 34%. If we exclude three

SNRs (G54.1+0.3, G116.9+0.2, and G184.6-5.8) which appear to

be outliers, we derive

NH(H cm−2) = (2.51±0.10)×1021 AV (mag) (14)

for SNRs, and

NH(H cm−2) = (2.47±0.04)×1021 AV (mag) (15)

for the whole W00 sample. Similar to the AG89 sample, with the

three SNR outliers excluded, the 〈NH/AV 〉 of SNRs agree with

that of XBs. We tabulate the best-fit results in Table 5.

The 〈NH/AV 〉 derived here is smaller by ∼ 15% than that

of Foight et al. (2016) who assumed subsolar abundances for 17

SNRs (see Table 3) and appreciably larger than most of those de-

rived from HI and/or H2 UV absorption and HI 21 cm emission

(see Tables 1,2). Two factors might be responsible for the inconsis-

tency. First, our sample is larger than previous ones and it includes

29 XBs which do not suffer the problem with which SNRs may

be confronted (i.e., the measured NH and AV may not be for the

same region. Second, the NH column derived from the UV absorp-

tion data of HI and/or H2 and the HI 21 cm emission data only

counts neutral and/or molecular hydrogen, while the NH value de-

termined from the X-ray absorption data includes atomic, molecu-

lar and ionic hydrogen.

With 〈NH/AV 〉 ≈ 2.5×1021 mag cm2 H−1, we estimate the

Figure 9. Top (a): NH against AV for the whole W00 sample. Middle (b):

NH against AV for the XBs of the W00 sample. Bottom (c): NH against

AV for the SNRs of the W00 sample.

c© 0000 RAS, MNRAS 000, 000–000


Figure 10. NH derived from subsolar abundances (W00) against NH de-

rived from solar abundances (AG89).

gas-to-dust mass ratio to be Mgas/Mdust ≈ 170, again, based on

the mass extinction coefficient of Li & Draine (2001) and Draine &

Li (2007). As expected, the gas-to-dust mass ratio of the W00 sam-

ple exceeds that of the AG89 sample by ∼ 20%. In Figure 10, we

plot NH from the W00 sample against that from the AG89 sample.

The fact that they are closely related with NH(W00)/NH(W00) ≈1.2 suggests that in future studies one may simply scale NH derived

from subsolar abundances by a factor of ∼ 1.2 when converting to

NH derived from solar abundances.

3.6 Caveats

Even though in our analysis we have tried our best to avoid system-

atic uncertainties (e.g., for SNRs adopting NH and AV determined

from the same region if possible, excluding those XBs with intrin-

sic local X-ray absorptions), there may still be sources of potential

uncertainty:

• Since we compile NH and AV from the literature, for some

SNRs the adopted NH and AV may not be for the same region

(e.g., SNR G132.7).

• Since a number of SNRs only have one or two measures of NH

and/or AV , we are not able to take into account the uncertainties

caused by spatial variations.

• For most SNRs, AV was measured from line ratios, e.g.,

Hα(6563 A)/Hβ(4861 A), [SII](∼10320 A)/[SII](∼4068 A), and

[FII](∼16435 A)/[FII](∼12567 A). Taking the Balmer decrement

as an example, the intrinsic ratio of Hα/Hβ is commonly taken to

be ∼ 3 under the “Case A” condition, in which the nebula is always

optically thin even for Lyman-line photons. However, this may not

be always true for SNRs. Previous studies have showed that Hα/Hβcould vary from ∼ 2.9 to ∼ 4.2 and even up to ∼ 6 (e.g., see Ray-

mond 1979, Shull & McKee 1979, Ghavamian et al. 2001).

• For some XBs, the optical extinction was estimated by com-

paring the observed colour with the intrinsic colour derived from

the stellar spectral type. However, if the optical flux contribution of

the accretion disk or the absorption of the stellar wind is not negli-

gible, the derived spectral type of the primary star may be incorrect

and this may result in an incorrect AV (e.g., for EXO 0748-676,

Schoembs & Zoeschinger (1990) estimated AV ≈ 1.3 ± 0.1magbased on its stellar spectral type, while Pearson et al. (2006) es-

timated AV ≈ 0.48 ± 0.26mag from the widths of the sodium

doublet absorption lines).

4 SUMMARY

The NH–AV relation is explored based on the X-ray absorption

data of a large AG89 sample of Galactic lines of sight toward 35 su-

pernova remnants, 6 planetary nebulae, 70 X-ray binaries for which

NH was derived with solar abundances, as well as a smaller W00

sample of 19 supernova remnants and 29 X-ray binaries for which

NH was derived with subsolar abundances. For the AG89 sample,

our principal results are as follows:

(i) We derive a mean hydrogen-to-extinction ratio of

〈NH/AV 〉 = (2.08± 0.02) × 1021 Hcm−2 mag−1 for the

whole Galaxy.

(ii) The X-ray-derived NH/AV ratio is generally constant

across the Galaxy. It does not appear to correlate with nH

(the line of sight mean number density of hydrogen), Rg (the

distance away from the Galactic centre), and z (the distance

above or below the Galactic plane). The 〈NH/AV 〉 ratio de-

rived for the 1st and 4th galactic quadrants (〈NH/AV 〉 =(2.04± 0.05) × 1021 Hcm−2 mag−1) is consistent with that of

the 2nd and 3rd galactic quadrants (〈NH/AV 〉 = (2.09± 0.03) ×1021 Hcm−2 mag−1).

(iii) We estimate a gas-to-dust mass ratio of Mgas/Mdust ≈140 from 〈NH/AV 〉 = 2.08 × 1021 Hcm−2 mag−1.

(iv) The distribution of hydrogen in the Galaxy is explored with

the addition of an additional sample of 74 supernova remnants for

which both NH and distances are known. It is found that, between

the Galactic radius of 2 kpc to 10 kpc, the vertical distribution of

hydrogen can be roughly described by a Gaussian function with a

mid-plane density of nH(0) = 1.11± 0.15 cm−3 and a scale hight

of h = 75.5±12.4 pc. We also estimate a total gas surface density

of∑

gas ∼ 7.0M⊙ pc−2.

Similarly, for the W00 sample we derive NH/AV =(2.47± 0.04)×1021 Hcm−2 mag−1 for the whole Galaxy which

is higher by ∼ 20% than that of the AG89 sample. We also find

that NH derived from subsolar abundances exceeds that from solar

abundances by ∼ 20%, suggesting that in future studies one may

simply scale NH of subsolar abundances by a factor of ∼ 1.2 when

converting to NH of solar abundances.

ACKNOWLEDGEMENTS

HZ and WWT acknowledge support from NSFC (10273025,

11473038, 11603039). AL is supported in part by NSF AST-

1311804 and NASA NNX14AF68G. We thank Patrick Slane, Linli

Yan, and the referee Tolga Guver for their very helpful comments

and suggestions which substantially improve the quality of our pa-

per. We also thank H. Su, S. Yao and D. Wu for their help during

the preparation of this paper.

APPENDIX: A LIST OF NH, AV AND D DATA COMPILED

FROM THE LITERATURE

c© 0000 RAS, MNRAS 000, 000–000

14

Zh

u,

Tia

n,L

i&

Zh

an

g

Table 7. Hydrogen column density NH and distance D compiled from the literature.

Object Type D References NH(AG89) References Accept “A” NH(AG89) NH(W00) References Accept “A” NH(W00) AV References Accept “A” AV(mag)

(kpc) for D (1021 H cm−2 ) for NH(AG89) or Reject

“R”

This Work (1021 H cm−2) for NH(W00) or Reject

“R”

This Work (mag) for AV or Reject

“R”

This Work

G4.5+6.8 SNR 5.1 Sankrit et al. 2016 5.3±0.4 Bamba et al. 2005 A 5.3±0.1 5.4±0.1 Foight et al. 2016 A 5.4±0.1 2.8±0.3 Blair et al. 1991 A 3.0±0.4

north 3.5±0.2 Cassam-Chenai et al. 2004 A 2.5±0.9 Oliva et al. 1989 A

5.7±0.1 Kinugasa & Tsunemi 1999 A 3.5 Leibowitz & Danziger 1983 A

3.3 Dennefeld 1982 A

2.2±0.6 van den Bergh & Kamper 1977 A

G6.4-0.1 SNR 1.9 Velazquez et al. 2002 4.1±0.1 Zhou et al. 2014 A 4.3±0.1 6.3±1.9 Foight et al. 2016 A 6.3±1.9 3.4±0.3 Mavromatakis et al. 2004a A 3.3±0.3

Central 4.7±0.2 Sawada & Koyama 2012 A 3.1 Long et al. 1991 A

5.1±0.5 Kawasaki et al. 2005 A

6.6±0.6 Rho & Borkowski 2002 A

G6.4-0.1 SNR 1.9 Velazquez et al. 2002 7.9±0.2 Zhou et al. 2014 A 7.7±0.2 4.0 Long et al. 1991 A 4.0±0.5

northeast shell 7.6±0.2 Nakamura et al. 2014 A

7.7±0.8 Rho & Borkowski 2002 A

G11.2-0.3 SNR 4.4 Kilpatrick et al. 2016 17.1±3.3 Roberts et al. 2003 A 15.7±1.5 13.0 Koo et al. 2007 A 13.0±1.7

southeast 15.9±1.1 Vasisht et al. 1996 A

13±3 Reynolds et al. 1994 A

G13.3-1.3 SNR 3.0 Seward et al. 1995a 5.5±4.5 Seward et al. 1995a A 5.5±4.5 0.5 Seward et al. 1995a A 0.5±0.1

G39.7-2.0 SNR 5.5 Blundell & Bowler 2004 5.2±0.1 Brinkmann et al. 2007 A 5.2±0.1 2.3±0.3 Boumis et al. 2007 A 2.3±0.3

east jet 2.2±0.6 Brinkmann et al. 1996 R

3.3±0.6 Yamauchi et al. 1994 R

G53.6-2.2 SNR 2.3 Giacani et al. 1998 6.1±0.1 Broersen & Vink 2015 A 6.1±0.1 3.1 Long et al. 1991 A 3.1±0.4

central 6.2±1.1 Yoshita et al. 2001 A

G53.6-2.2 SNR 2.3 Giacani et al. 1998 4.8±0.7 Broersen & Vink 2015 A 4.8±0.8 6.7±0.3 Foight et al. 2016 A 6.7±0.3 4.0 Long et al. 1991 A 4.0±0.5

southwest 4.8±1.1 Yoshita et al. 2001 A

G54.1+0.3 SNR 6.2 Leahy et al. 2008 19.5±0.4 Temim et al. 2010 A 18.1±0.5 25.5±0.4 Foight et al. 2016 A 25.5±0.4 7.3±0.1 Kim et al. 2013 A 7.3±0.1

15.7±0.6 Bocchino et al. 2010 A 8.0±0.7 Koo et al. 2008 A

16±1 Lu et al. 2002 A

17.9±2.7 Lu et al. 2001 A

G67.7+1.8 SNR 4.1±0.9 Hui & Becker 2009 A 4.1±0.9 1.0-2.5 Gok et al. 2008 A 1.7±0.4

north rim 1.7±0.3 Mavromatakis et al. 2001 A

G69.0+2.7 SNR 1.5 Leahy & Ranasinghe 2012 3.0±0.1 Li et al. 2005 A 3.0±0.1 4.5±0.2 Foight et al. 2016 A 4.5±0.2 2.4 Hester & Kulkarni 1989 A 2.5±0.3

central 2.9±0.5 Mavromatakis et al. 2001 A 2.4 Blair et al. 1984a A

3.1-4.3 Angerhofer et al. 1980 A

G74.0-8.5 SNR 1.5 Blair et al. 2005 0.95±0.13 Oakley et al. 2013 R 0.34±0.03 0.25 Miller 1974 A 0.25±0.04

0.34±0.02 Leahy & Hassan 2013 A 0.25±0.06 Parker 1967 A

0.35±0.05 Kimura et al. 2009 A

0.01-1.47 Nemes et al. 2008 R

0.4±0.1 Uchida et al. 2008 A

0.5±0.1 Miyata et al. 2007 A

G78.2+2.1 SNR 2.0 Leahy et al. 2013 7.4±0.6 Hui et al. 2015 A 7.6±0.7 3.4±0.6 Mavromatakis 2003 A 3.4±0.6

east 12±2 Uchiyama et al. 2002 R

9.2±1.8 Leahy et al. 2013 A

for central

G82.2+5.3 SNR 2.0 Uyaniker et al. 2003 4.7±0.5 Mavromatakis et al. 2004b A 4.7±0.5 2.8±0.1 Mavromatakis et al. 2004b A 2.8±0.1

G85.9-0.6 SNR 5.0 Kothes et al. 2001 6.9±0.3 Jackson et al. 2008 A 6.9±0.3 0.7±0.1 Gok et al. 2009 A 0.7±0.1

G89.0+4.7 SNR 1.7 Byun et al. 2006 2.6±0.5 Pannuti et al. 2010b A 3.0±0.4 1.6±0.3a Mavromatakis et al. 2007 A 1.6±0.3

3.2±0.3 Lazendic & Slane 2006 A

G109.1-1.0 SNR 3.2 Kothes & Foster 2012 5-7 Sasaki et al. 2004 A 7.1±0.7 9.0±1.4 Foight et al. 2016 A 9.0±1.4 3.1±0.6 Fesen & Hurford 1995 A 3.1±0.6

6.9±0.8 Parmar et al. 1998 A

7.7±0.6 Rho & Petre 1997 A

G111.7-2.1 SNR 3.4 Reed et al. 1995 12.7±0.3 Laming et al. 2006 A 13.6±0.3 11.7±1.6 Foight et al. 2016 A 11.7±1.6 4.6-6.2 Hurford & Fesen 1996 A 4.8±0.7

north 14.2±0.2 Laming & Hwang 2003 A 4.4 Searle 1971 A

15±8 Willingale et al. 2002 A

10±1 Gotthelf et al. 2001 A

G116.9+0.2 SNR 3.0 Yar-Uyaniker et al. 2004 6.3±1.2 Pannuti et al. 2010b A 5.9±0.8 9.2±0.7 Foight et al. 2016 A 9.2±0.7 2.6±0.1 Fesen et al. 1997 A 2.6±0.1

5.7±0.7 Lazendic & Slane 2006 A

c©0000

RA

S,M

NR

AS

000,

000–000

Th

eG

ala

cticN

H–A

VR

elatio

n1

5

G119.5+10.2 SNR 1.4 Pineault et al. 1993 1.8±0.3 Caraveo et al. 2010 A 2.1±0.4 3.8±1.1 Foight et al. 2016 A 3.8±1.1 1.27±0.14 Mavromatakis et al. 2000 A 1.3±0.2

2.8±0.6 Slane et al. 1997 A 1.5 Fesen et al. 1981 A

2.3±0.6 Seward et al. 1995b A

G120.1+1.4 SNR 3.0 Zhou et al. 2016 6.6±0.4 Cassam-Chenai et al. 2007 A 5.8±0.4 8.0±1.0 Foight et al. 2016 A 8.0±1.0 1.9±0.1 Ruiz-Lapuente 2004 A 1.8±0.1

Tian & Leahy 2011 6.4±0.5 Warren et al. 2005 A ∼1.6 Ghavamian et al. 2000 A

4.7±0.3 Hwang et al. 2002 A 2.1±0.5 Chevalier et al. 1980 A

8.3 Decourchelle et al. 2001 A

G130.7+3.1 SNR 2.0 Kothes 2013 4.2±0.1 Gotthelf et al. 2007 A 4.1±0.1 5.4±0.1 Foight et al. 2016 A 5.4±0.1 2.1±0.2 Fesen et al. 1988 A 2.1±0.2

4.5±0.1 Slane et al. 2004 A 1.9±0.9 Fesen 1983 A

3.8±0.1 Slane et al. 2002 A

3.3±0.4 Torii et al. 2000 A

3.3±0.7 Helfand et al. 1995 A

G132.7+1.3 SNR 2.5 Routledge et al. 1991; Landecker

et al. 1987

8.6±0.3 Lazendic & Slane 2006 A 8.6±0.3 2.2±0.1 Fesen et al. 1995 A 2.2±0.1

G166.0+4.3 SNR 4.5 Landecker et al. 1989 ∼1.9 Bocchino et al. 2009 A 2.0±0.2 1.7 Fesen et al. 1985 A 1.7±0.2

∼2.1 Guo & Burrows 1997 A

±2 Burrows & Guo 1994 A

G180.0-1.7 SNR 1.3 Chatterjee et al. 2009 2.9±0.1 Ng & Romani 2007 A 3.1±0.2 0.7±0.2 Fesen et al. 1985 A 0.7±0.2

3.9±0.2 Romani & Ng 2003 A

G184.6-5.8 SNR 2.0 Trimble 1968 ∼3.5 Seward et al. 2006 A 3.2±0.1 3.0±0.2 Foight et al. 2016 A 4.1±0.1 1.6±0.2 Miller 1973, Sollerman et al.

2000

A 1.5±0.1

2.7±0.1 Kirsch et al. 2006 R 4.4±0.1 Weisskopf et al. 2004 A 1.6 Sollerman et al. 2000 A

3.6±0.1 Weisskopf et al. 2004 A 1.4±0.1 Veron-Cetty & Woltjer 1993 A

3.5±0.1 Willingale et al. 2001 A 1.6±0.1 Blair et al. 1992 A

1.6±0.1 Wu 1981 A

G189.1+3.0 SNR 1.5 Frail 2011 ∼4.1 Bocchino et al. 2009 A 5.2±0.5 ∼6 Yamaguchi et al. 2014 A 6±0.5 ∼2-5 Kokusho et al. 2013 A 2.8±0.6

5.2±0.4 Troja et al. 2006 A 2.6±0.7 Welsh & Sallmen 2003 A

7.2±0.6 Gaensler et al. 2006 A 2.9±0.5 Fesen & Kirshner 1980 A

8.0±1.5 Bocchino & Bykov 2001 A

G260.4-3.4 SNR 2.0 Reynoso et al. 1995 2.7±0.1 Katsuda et al. 2012 A 3.1±0.1 3.5±1.5 Foight et al. 2016 A 3.5±1.5 1.7±0.6 Blair et al. 1995 A 1.5±0.2

Sakhibov & Smirnov 1983 4.3±0.2 Katsuda et al. 2008 A 1.5 Winkler & Kirshner 1985 A

4.5±0.5 Hui & Becker 2006 A 0.6 Dopita et al. 1977 R

G263.9-3.3 SNR ∼0.29 Verbiest et al. 2012 0.13±0.09 Yamaguchi & Katsuda 2009 A 0.19±0.02 0.3±0.1 Foight et al. 2016 A 0.3±0.1 0.31 Sankrit et al. 2003 A 0.32±0.04

0.23±0.09 Miceli et al. 2008 A 0.1 Bocchino et al. 2000 R

0.16±0.03 Lamassa et al. 2008 A 0.56 Raymond et al. 1997 R

0.14±0.02 Katsuda & Tsunemi 2006 A 0.31 Wallerstein & Balick 1990 A

0.24±0.03 Katsuda & Tsunemi 2005 A

0.15±0.25 Miceli et al. 2005 A 0.37 Manchester et al. 1978 A

0.19±0.1 Lu & Aschenbach 2000 A

0.03-0.3 Bocchino et al. 1994 A

0.35±0.08 Tsunemi et al. 1999 A

G292.0+1.8 SNR 6.0 Gaensler et al. 2003 4.3±0.2 Kamitsukasa et al. 2014 A 3.8±0.3 ∼5 Yamaguchi et al. 2014 A 5±0.4 1.9 Winkler & Long 2006 A 2.2±0.2

5.4±0.6 Lee et al. 2010 A 2.8±0.3 Goss et al. 1979 A

5.2±0.6 Park et al. 2004 A

5±1 Gonzalez & Safi-Harb 2003 A

3.1±0.4 Hughes et al. 2003 A

3.2±0.2 Hughes et al. 2001 A

G296.1-0.5 SNR ∼2.0 Castro et al. 2011 0.75±0.05 Gok & Sezer 2012 A 0.6±0.2 1.9±0.3 Longmore et al. 1977 A 1.9±0.3

north 0.2±0.1 Castro et al. 2011 A

G296.5+10.0 SNR ∼2.1 Giacani et al. 2000 1.2±0.1 De Luca et al. 2004 A 1.2±0.1 0.5 Ruiz 1983 A 0.5±0.1

1.5±0.4 Zavlin et al. 1998 A

G315.4-2.3 SNR 2.3 Sollerman et al. 2003 5.1±0.4 Helder et al. 2011 A 4.0±0.2 1.7 Leibowitz & Danziger 1983 A 1.7±0.2

southwest 4.3±0.6 Bamba et al. 2005 A

4.3±0.2 Rho et al. 2002 A

2.7±1.0 Borkowski et al. 2001 A

3.2±0.5 Bamba et al. 2000 A

∼2.8 Bocchino et al. 2000 A

2.0±0.4 Vink et al. 1997 A

1.8±0.4 Lemoine-Goumard et al. 2012 A

4.8±0.1 Broersen et al. 2014 A

G320.4-1.2 SNR 5.2 Gaensler et al. 1999 8.6±0.2 DeLaney et al. 2006 A 10.9±0.1 11.5±0.3 Schock et al. 2010 A 11.5±0.3 5.5 Seward et al. 1983 A 3.6±0.4

11.7±0.1 Yatsu et al. 2005 A 3.0-3.6 Dopita et al. 1977 A

9.5±0.3 Gaensler et al. 2002 A

12.7±2.3 Marsden et al. 1997 A

5.9±0.6 Tamura et al. 1996 A

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G327.6+14.6 SNR 2.2 Winkler et al. 2003 0.68±0.01 Uchida et al. 2013 A 0.68±0.10 1.6±0.1 Foight et al. 2016 A 1.6±0.1 0.32±0.1 Schweizer & Middleditch 1980 A 0.32±0.1

0.42±0.08 Broersen et al. 2013 A

0.90±0.05 Bamba et al. 2003 A

0.56±0.06 Allen et al. 2001 A

G332.4-0.4 SNR 3.1 Caswell et al. 1975 7 Gotthelf et al. 1997 A 7±0.7 8.7±3.5 Foight et al. 2016 A 8.7±3.5 4.5 Leibowitz & Danziger 1983 A 4.7±0.7

A ∼5.0 Ruiz 1983 A

4.7±0.9 Oliva et al. 1989 A

G332.5-5.6 SNR 3.0 Zhu et al. 2015 1.4±0.4 Zhu et al. 2015 A 1.4±0.4 0.83 Stupar et al. 2007 A 0.8±0.3

0.83 Stupar et al. 2007 R

G37.7-34.5 PN 1.1 Gaensler et al. 1999 0.81±0.85 Guerrero et al. 2002 A 0.81±0.85 0.29±0.23 Rubin et al. 2001 A 0.19±0.02

0.70±0.18 Lame & Pogge 1996 R

0.19 Mendez et al. 1988 A

G64.7+5.0 PN 5.2 Gaensler et al. 1999 2.5±0.1 Maness et al. 2003 A 2.5±0.1 1.1 Bernard-Salas et al. 2003 A 1.0±0.1

0.9 Aller & Hyung 1995 A

0.9 Rudy et al. 1991 A

0.9 Torres-Peimbert & Pena 1981 A

G84.9-3.4 PN 0.7 Yang et al. 2016 4.1±0.2 Maness et al. 2003 A 4.1±0.2 2.6 Salas et al. 2001 A 2.9±0.3

2.9±0.2 Aller & Keyes 1988 A

2.9 Pottasch et al. 1982 A

3.0 Shaw & Kaler 1982 A

3.0±0.3 Atherton et al. 1979 A

G94.0+27.4 PN 1.0 Cahn et al. 1992 0.12±0.02 Montez & Kastner 2013 A 0.12±0.02 0.09±0.03 Kaler & Feibelman 1985 A 0.09±0.03

G197.8+17.3 PN 1.2 Bensby & Lundstrom 2001 0.8±0.5 Montez et al. 2015 A 0.8±0.6 0.5±0.1 Pottasch et al. 2008 A 0.5±0.1

0.8±0.9 Guerrero et al. 2005 A

G331.7-1.0 PN 1.1 Ortiz et al. 2011 5.0±0.8 Kastner et al. 2003 A 5.0±0.8 3.2 Smith 2003 A 4.1±0.4

4.6±0.3 Zhang & Liu 2002 A

4.4 Tylenda et al. 1992 A

2E 0236.6+6101 HMXB 2.4 Liu et al. 2006 6.1±0.1 Rea et al. 2010 A 5.9±0.2 3.5 Paredes & Figueras 1986 A 2.8±0.3

5.1±0.4 Albert et al. 2008 A 2.3 Howarth 1983 A

5.7±0.3 Esposito et al. 2007 A 2.8±0.2 Maraschi et al. 1981 A

7.0±0.6 Paredes et al. 2007 A 2.9 Hutchings & Crampton 1981a A

4.9±0.2 Chernyakova et al. 2006a A

6±1 Sidoli et al. 2006 A

6.0±0.5 Leahy et al. 1997 A

2E 1118.7-6138 HMXB 5.0 Liu et al. 2006 8.1±0.1 Maitra et al. 2012 A 8.1±0.1 13.2±0.3 Suchy et al. 2011 A 13.2±0.3 2.9±0.4 Coe et al. 1994a A 2.9±0.3

2.8±0.2 Coe & Payne 1985 A

3.7±0.6 Janot-Pacheco et al. 1981 A

2S 0620-003 LMXB 1.2 Liu et al. 2007 2.4±1.1 Kong et al. 2002 A 2.4±1.1 1.1±0.1 Wu et al. 1983 A 1.2±0.1

1.2 Greenstein 1979 A

1.2±0.1 Maraschi et al. 1981 A

2S 0921-63 LMXB 8.5 Liu et al. 2007 2.3±0.1 Valencic & Smith 2015 A 2.3±0.1 1.02±0.03 Shahbaz & Watson 2007 A 1.02±0.03

4U 0115+63 HMXB 17.5±12.3 Li et al. 2012 R 5.1±0.2 Negueruela et al. 2001 A 5.1±0.2

17.4±1.8 Campana et al. 2001 R 4.7 Hutchings & Crampton 1981b A

4U 0352+30 HMXB 1.0 Liu et al. 2006 2.5±0.1 Valencic & Smith 2015 A 2.5±0.1 1.2 Telting et al. 1998 A 1.2±0.1

3.4±0.2 La Palombara & Mereghetti 2007 A 1.2±0.1 Fabregat et al. 1992 A

1.5±0.2 Di Salvo et al. 1998 A 1.1 Bernacca et al. 1983 A

1.2±0.1 Ferrari-Toniolo et al. 1978 A

4U 0538+26 HMXB 2.5 Liu et al. 2006 3.5±0.1 Doroshenko et al. 2014 A 3.9±0.1 7.0±0.3 Caballero et al. 2013 A 7.0±0.3 2.3 Lyuty & Zaitseva 2000 A 2.4±0.1

5.3±0.2 Maitra & Paul 2013 A 2.2±0.1 Steele et al. 1998 A

6.7±1.5 Mukherjee & Paul 2005 A 2.3 De Loore et al. 1984 A

2.2 Wu et al. 1983 A

2.5±0.1 Giangrande et al. 1980 A

2.5 Wade & Oke 1977 A

2.5 Stier & Liller 1976 A

4U 0614+09 LMXB 3.4±0.3 Migliari et al. 2010 A 3.5±0.1 3.3±0.1 Schulz et al. 2010 A 3.3±0.1 2.7±0.3 Nelemans et al. 2004 A 1.2±0.1

1.5-2.0 Paerels et al. 2001 R 1.0 Davidsen et al. 1974 A

3.8±0.1 Juett et al. 2001 A

3.2±0.1 Piraino et al. 1999 A

4U 0919-54 LMXB 2.8±0.1 Valencic & Smith 2015 A 3.1±0.1 3.0±0.1 Juett & Chakrabarty 2003 A 3.0±0.1 2.2±0.5 Nelemans et al. 2004 R 1.0±0.1

3.0±0.3 in’t Zand et al. 2005 A 1.0±0.3 Chevalier & Ilovaisky 1987 A

3.5±0.1 Juett et al. 2001 A 1.2 Chevalier & Ilpvaisky 1985 A

1.0 Neckel et al. 1980 A

4U 1036-56 HMXB 5.0 Liu et al. 2007 22±8 Cusumano et al. 2013 R 2.8±0.9 2.3±0.8 Motch et al. 1997 A 2.3±0.8

>14 Li et al. 2012 R

2.8±0.9 La Palombara et al. 2009 A

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4U 1118-60 HMXB 9.0 Liu et al. 2006 11.2±0.1 Naik et al. 2011a A 11.0±0.1 4.3 van der Meer et al. 2007 A 4.0±0.5

9.5±0.3 Wojdowski et al. 2001 A

2.8 Coe & Payne 1985 R

3.7±0.6 Janot-Pacheco et al. 1981 A

4.2 Hutchings et al. 1979 A

3.3 Genderen 1975 A

4.2 Mauder 1975 A

4.4 Krzeminski 1974 A

4.2 Vidal et al. 1974 A

4U 1258-61 HMXB 2.4 Liu et al. 2007 10.0±0.2 Jaisawal & Naik 2015 A 10.0±0.2 5.3±0.6 Parkes et al. 1980 5.5±0.7

5.9 Menzies et al. 1981

4U 1254-69 LMXB 13.0 Liu et al. 2007 2.0±0.1 Iaria et al. 2007 A 2.6±0.1 2.2±0.1 Trigo et al. 2008 A 2.2±0.1 1.2±0.3 Motch et al. 1987 A 1.2±0.3

3.2±0.1 Boirin & Parmer 2003 A

3.5±0.5 Iaria et al. 2001 A

4U 1456-32 LMXB 1.2 Liu et al. 2007 0.5±0.1 Cackett et al. 2010 A 0.7±0.1 0.9±0.1 Chakrabarty et al. 2014 A 0.9±0.1 0.31±0.15 Blair et al. 1984b A 0.31±0.15

3±1 Asai et al. 1998 R 0.8±0.4 Campana et al. 2004 A

0.8±0.1 Bernardini et al. 2013 A

4U 1516-56 HMXB 5.5 Liu et al. 2007 18±1 Heinz et al. 2013 A 18.4±1.2 12.2 Jonker et al. 2007 A 12.2±1.6

20.3±2.3 Sell et al. 2010 A 12.1 Whelan et al. 1977 A

4U 1538-52 HMXB 5.5 Liu et al. 2006 14.3±2.6 Valencic & Smith 2015 A 16.2±0.5 20.6±0.5 Hemphill et al. 2016 A 20.6±0.5 6.9 Pakull et al. 1983 A 7.1±0.7

46 Clark 2004 R 6.8 Ilovaisky et al. 1979 A

16.3±0.4 Robba et al. 2001 A 7.4±0.5 Parkes et al. 1978 A

6.8 Crampton et al. 1978 A

4U 1543-47 LMXB 7.5 Liu et al. 2007 3.8±0.2 La Palombara & Mereghetti 2005 A 3.8±0.2 1.6±0.2 Orosz et al. 1998 A 1.6±0.2

A

4U 1543-62 LMXB 7.0 Liu et al. 2007 2.6±0.9 Farinelli et al. 2003 A 3.2±0.1 3.4±0.1 Juett & Chakrabarty 2003 A 3.4±0.1 1.0±0.3 Nelemans et al. 2004 A 1.0±0.3

3.3±0.7 Schultz 2003 A

1.4±0.3 Asai et al. 2000 A

3.4±0.1 Juett et al. 2001 A

4U 1608-52 LMXB 3.9 Liu et al. 2007 13.6±0.4 Valencic & Smith 2015 A 10.7±0.2 10.8±2.2 Guver et al. 2010a A 10.8±2.2 >2 Wachter et al. 2002 R 6.2±1.7

9±6 Asai et al. 1996 A 6.2±1.7 Grindlay & Liller 1978 A

10.5±0.1 Armas Padilla et al. 2017 A

4U 1636-53 LMXB 6.0 Liu et al. 2007 4.4±1.4 Cackett et al. 2009a A 2.9±0.1 3.5±0.6 Sanna et al. 2013 A 3.7±0.1 2.5±0.3 Lawrence et al. 1983 A 2.5±0.3

2.9±0.1 Lyu et al. 2014 A 3.7±0.2 Pandel et al. 2008 A

3.7±0.1 Lyu et al. 2015 A

4U 1658-48 LMXB 2.9±0.1 Plant et al. 2015 R 5.0±0.1 7.7±0.2 Parker et al. 2016 A 8.1±0.2 3.3±0.5 Gandhi et al. 2011 A 3.7±0.4

5.4±0.1 Shidatsu et al. 2011b A 8.5±0.2 Furst et al. 2015 A 3.9 Casella et al. 2010 A

5.3±0.1 Miller et al. 2008 A 6.5±1.2 Garcia et al. 2015 A 3.4±0.6 Buxton & Vennes 2003 A

5.2±0.1 Reis et al. 2008 A 3.7±0.3 Zdziarski et al. 1998 A

4.0±0.1 Miller et al. 2004a A 4.0 Cowley et al. 1987 A

5.3±0.1 Miller et al. 2004b A 2.2±0.3 Ilovaisky et al. 1986 R

5.1±1.5 Kong et al. 2000a A 3.9±0.8 Grindlay 1979 A

4.5±0.1 Shidatsu et al. 2011a

4U 1704-30 LMXB 10.0 Liu et al. 2006 2.0±0.2 Cackett et al. 2006 A 2.3±0.2 0.9 Wachter & Smale 1998 A 0.9±0.1

2.4±0.8 Wijnands et al. 2003 A

3.5±0.2 Sidoli et al. 2001a A

1.3±0.2 Oosterbroek et al. 2001a A

4U 1724-307 LMXB 7.4 Liu et al. 2007 6.6±0.5 Valencic & Smith 2015 A 6.6±0.5 5.8±0.6 Harris 2010 and 1996 A 4.9±0.6

6.5±0.5 Valencic et al. 2009 A 5.8 Valenti et al. 2010 A

11.0±0.7 Barret et al. 1999 R 4.8 rtolani et al. 1997 A

3.9±0.5 Christian & Friel 1992 A

4U 1728-16 LMXB 4.2b Green et al. 2014 2.4±0.2 Xiang et al. 2005 A 2.4±0.1 0.8±0.5 Shahbaz et al. 1996 A 0.8±0.5

2.4±0.1 Church & Balucinska-Church

2001

A

4U 1728-34 LMXB 5.2 Liu et al. 2007 23.4±1.0 Egron et al. 2011 A 29.3±0.5 7.8±1.7 Di Salvo et al. 2000 A 10.8±1.8

25.1±0.8 D’Ai et al. 2006 A 15.0 Glass 1978 A

31.3±0.3 Narita et al. 2001 A

28.6±1.0 Di Salvo et al. 2000 A

27.3±0.5 Piraino et al. 2000 A

22.9±2.5 Misanovic et al. 2010 A

4U 1730-335 LMXB ∼8 Liu et al. 2007 17.4±0.3 Marshall et al. 2001 A 17.4±0.3 9.5 Harris 2010 and 1996 A 10.1±1.1

9.3 Homer et al. 2001 A

11±1 Kleinmann et al. 1976 A

4U 1746-37 LMXB 11.0 Liu et al. 2007 4.0±0.3 Gavriil et al. 2012 A 4.0±0.3 1.5±0.3 Bonatto et al. 2013 A 1.5±0.2

2.8±0.4 Parmar et al. 1999 R 1.5±0.2 Harris 2010 and 1996 A

1.6 Deutsch et al. 1998 A

1.4±0.5 Hesser & Hartwick 1976 A

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4U 1820-30 LMXB 7.6 Liu et al. 2007 2.9±0.1 Valencic & Smith 2015 A 2.7±0.1 2.5±0.3 Guver et al. 2010b A 2.6±0.3 0.9±0.1 Harris 2010 and 1996 A 0.9±0.1

2.0 Yao & Wang 2006 A 1.9±0.1 Mondal et al. 2016 R 1.1 Valenti et al. 2004 A

2.7±0.4 Juett et al. 2004 A 1.0 Heasley et al. 2000 A

0.9 Reed et al. 1988 A

0.8 Zinn 1980 A

4U 1822-37 LMXB 2.5 Liu et al. 2007 1.2±0.2 Iaria et al. 2001 A 1.2±0.2 0.4±0.4 Mason et al. 1982 A 0.5±0.3

0.8±0.7 Heinz & Nowak 2001 A 0.6±0.3 Mason & Cordova 1982 A

4U 1837+04 LMXB 4.9±0.2 Xiang et al. 2005 A 5.2±0.2 8.9±0.1 Matranga et al. 2017 A 2.5±1.2 Hynes et al. 2004 A 1.3±0.1

5.0±0.3 Oosterbroek et al. 2001b A 1.3±0.1 Shahbaz et al. 1996 A

6.5±1.4 Seon & Min 2002 A >1.2 Thorstensen, Charles & Bowyer

1980

A

4U 1850-08 LMXB 8.2 Liu et al. 2007 4.0±0.1 Valencic & Smith 2015 A 3.8±0.1 4.8±0.2 Sidoli et al. 2005a A 4.8±0.2 1.4±0.2 Harris 2010 and 1996 A 1.2±0.2

4±2 Sidoli et al. 2006 A 1.0± 0.2 Paltrinieri et al. 2001 A

3.9±0.3 Sidoli et al. 2001b A 1.0 Ortolani et al. 2000 A

2.9±0.2 Juett et al. 2001 A 1.4 Peterson 1993 A

1.3 Cudworth 1988 A

1.5 Webbink 1985 and Zinn 1985 A

1.2 Zinn 1980 A

4U 1908+00 LMXB 5.0 Liu et al. 2007 5.0±0.1 Valencic & Smith 2015 A 4.4±0.1 4.0±0.1 Campana et al. 2014 A 4.1±0.1 1.6±0.3 Chevalier et al. 1999 A 1.2±0.1

3.8±0.1 Cackett et al. 2011 A 5.9±0.4 Coti Zelati et al. 2014 A 1.1 Shahbaz et al. 1996 A

3.4±0.7 Campana & Stella 2003 A 1.3±0.2 Thorstensen et al. 1978 A

4U 1915-05 LMXB 8.9 Liu et al. 2007 4.4±0.1 Iaria et al. 2006 A 3.9±0.1 5.6±0.7 Juett & Chakrabarty 2006 5.6±0.7 2.2±0.9 Schmidtke 1988 A 2.2±0.9

5.8±0.2 Boirin et al. 2004 A

3.0±0.1 Ko et al. 1999 A

3.2±0.2 Church et al. 1998 A

4U 1956+35 HMXB 2.14 Liu et al. 2006 4.4±0.2 Xiang et al. 2005 A 4.4±0.1 6.4±0.2 Tomsick et al .2014 A 5.4±0.1 3.4±0.1 Caballero-Nieves et al. 2009 A 3.3±0.1

4.4±0.1 Cui et al. 1998 A 5.1±0.1 Parker et al. 2015 A 3.4 Ziolkowski 2005 A

5.4±0.1 Miskovicova et al. 2016 A 3.3 Bradt & McClintock 1983 A

2.9±0.2 Wu et al. 1982 A

4U 1957+11 LMXB 2.2±0.1 Valencic & Smith 2015 A 1.6±0.1 1.1±0.1 Nowak et al. 2008 A 1.1±0.1 0.9 Margon et al. 1978 A 0.9±0.1

1.2±0.1 Maitra et al. 2014 A

1.3±0.1 Nowak & Wilms 1999 A

4.4±2.9 Singh et al. 1994 R

4U 2129+47 LMXB 7.0 Liu et al. 2007 2.4±0.7 Bozzo et al. 2007 A 2.4±0.7 3.6±0.1 Lin et al. 2009 A 3.6±0.1 1.6±0.3 Shahbaz et al. 1996 A 1.2±0.2

3.2±0.6 Nowak et al. 2002 A 1.2 Deutsch et al. 1996 A

0.9 Cowley & Schmidtke 1990 A

1.6 Chevalier et al. 1989 A

1.2 Thorstensen et al. 1988 A

4U 2135+57 HMXB 3.8 Liu et al. 2006 7.8±0.3 Jaisawal et al. 2016 A 7.8±0.3 4.0±0.3 Bonnet-Bidaud & Mouchet 1998 A 4.0±0.3

4U 2142+38 LMXB 7.2 Liu et al. 2007 1.8±0.3 Seward & Smith 2013 A 1.9±0.1 3.6±0.1 Balucinska-Church et al. 2011 A 2.9±0.1 1.1 Shahbaz et al. 1996 A 1.1±0.1

2.2±0.4 Xiang et al. 2005 A 2.2±0.1 Costantini et al. 2005 A 0.7±0.2 Goranskij & Lyutyj 1988 A

2.0±0.1 Piraino et al. 2002 A 2.3±0.5 Juett et al. 2004 A 1.2±0.2 McClintock et al. 1984 A

1.9±0.1 Di Salvo et al. 2002 A 1.4±0.2 Chiappetti et al. 1983 A

EXO 0331+530 HMXB 16±6 Pottschmidt et al. 2005 A 9.0±1.0 5.9 Coe et al. 1987 A 5.8±0.2

9.1±1.3 Unger et al. 1992 A 5.8±0.3 Corbet et al. 1986 A

∼10-13 Tsygankov et al. 2016 A 6.8±0.8 Stocke et al. 1985 A

8.3±0.7 Caballero-Garcia et al. 2016 A 6.0±0.2 Honeycut & Schlegel 1985 A

10.7±4.3 Elshamouty et al. 2016 A 5.6±0.2 Kodaira et al. 1985 A

EXO 0748-676 LMXB 6.8 Liu et al. 2006 0.7±0.1 Degenaar et al. 2011 A 0.8±0.1 0.6±0.1 Degenaar et al. 2014 A 0.6±0.1 0.48±0.26 Pearson et al. 2006 A 0.2±0.1

1.3±0.3 Degenaar et al. 2009 A 0.6±0.1 Diaz Trigo et al. 2011 A 0.2±0.1 Hynes et al. 2006 A

1.1±0.4 Sidoli et al. 2005b A 1.3±0.1 Schoembs & Zoeschinger 1990 R

EXO 1745-248 LMXB 8.7 Liu et al. 2007 13±3 Degenaar & Wijnands 2012 A 13.2±3.5 7.0±1.0 Massari et al. 2012 A 7.2±1.0

14±5 Wijnands et al. 2005 A 7.1 Harris 2010 and 1996 A

7.4 Valenti et al. 2007 A

6.7 Cohn et al. 2002 A

7.4 Barbuy et al. 1998 A

7.7 Ortolani et al. 1996 A

5.1 Armandroff & Zinn 1988 R

EXO 2030+375 HMXB 7.1 Liu et al. 2006 19.9±0.2 Naik & Jaisawal 2015 A 20.5±0.1 29.8±0.3 Ferrigno et al. 2016 A 11.2±0.9 Janot-Pacheco et al. 1988 A 11.7±0.7

20.6±0.1 Naik et al. 2013 A 11.6 Coe et al. 1988 A

11.8±0.5 Motch & Janot-Pacheco 1987 A

GRO J1719-24 LMXB 1.2b Green et al. 2014 ∼4.0 Tanaka 1993 A 4±0.4 2.8±0.6 della Valle et al. 1994 A 2.8±0.6

GRO J0422+32 LMXB 2.49 Liu et al. 2007 1.6±0.9 Shrader et al. 1997 A 1.6±0.9 0.74±0.09 Gelino & Harrison 2003 A 0.8±0.1

0.9±0.3 Callanan et al. 1995 A

1.25±0.10 Chevalier & Ilovaisky 1995 A

0.72±0.05 Shrader et al. 1994 A

0.9±0.6 Shrader et al. 1993 A

0.9±0.3 Castro-Tirado et al. 1993 A

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GRO J1008-57 HMXB 5.0 Liu et al. 2006 9.6±0.5 Yamamoto et al. 2014 A 12.7±0.1 6.1±0.2 Coe et al. 1994b A 6.1±0.2

12.9±0.1 Naik et al. 2011b A

GRO J1655-40 LMXB 7.1±0.3 Kubota et al. 2005 A 7.4±0.3 7.1±0.3 Sala et al. 2007 A 7.4±0.2 3.7±0.2 Hynes et al. 1998 A 3.7±0.3

7.1±0.8 Yamaoka et al. 2001 A 7.7±0.2 Diaz Trigo 2007 A 3.5 Bianchini et al. 1997 A

7.6±0.2 Gierlinski et al. 2001 A 4.0 Horne et al. 1996 A

7.4±0.3 Ueda et al. 1998 A 3.6 Bailyn et al. 1995 A

GRO J1944+26 HMXB 9.5 Liu et al. 2007 13.0±0.1 Maitra & Paul 2013 A 13.0±0.1 16.3±0.3 Marcu-Cheatham et al. 2015 A 16.3±0.3 6.5±0.3 Verrecchia et al. 2002 A 6.5±0.3

GRO J2058+42 HMXB 9.0 Liu et al. 2007 21.0±10.0 Wilson et al. 2005 A 9.5±2.3 4.3±0.3 Wilson et al. 2005 A 4.3±0.4

9±2 Krimm et al. 2008 A 4.3 Reig et al. 2005a A

GRS 1009-45 LMXB 4.5 Liu et al. 2007 0.4±0.1 Kubota et al. 1998 A 0.4±0.1 0.6±0.2 della Valle et al. 1997 A 0.6±0.2

GS 0834-43 HMXB 4.0 Liu et al. 2006 24±3 Belloni et al. 1993 A 21.2±2.3 12.3 Israel et al. 2000 A 12.3±1.6

20±2 Aoki et al. 1992 A

GS 1124-684 LMXB 5.9 Liu et al. 2007 2.3±0.1 Greiner et al. 1994 A 2.3±0.1 0.93±0.16 Shrader & Gonzalez-Riestra 1993 A 0.90±0.1

0.90±0.03 Cheng et al. 1992 A

0.9±0.3 della Valle et al. 1991 A

HD 259440 HMXB 3.2±1.1 Rea & Torres 2011 R 2.6 Aragona et al. 2010 A 2.6±0.3

3.3±0.2 Falcone et al. 2010 R

3.1±0.3 Hinton et al. 2009 R

IGR

J11215-5952

HMXB ∼8.0 Liu et al. 2006 22.0±13.0 Romano et al. 2011 R 2.2 Lorenzo et al. 2014 A 2.3±0.3

10.4±2.5 Romano et al. 2007 R 2.4 Masetti et al. 2006 A

10.2±4.2 Mangano et al. 2007 R

12±3 Esposito et al. 2009 R

IGR

J17544-2619

HMXB 3.7b Marshall et al. 2006 11.2±0.8 Valencic & Smith 2015 A 11.2±1.2 6.3±0.4 Pellizza et al. 2006 A 6.3±0.4

12±2 Romano et al. 2015 A

14±3 Drave et al. 2014 A

11±2 Sidoli et al. 2009a A

18±10 Sidoli et al. 2009b A

19±5-

43±2

Gonzalez-Riestra et al. 2004 R

IGR

J18450-0435

HMXB 14.0b Marshall et al. 2006 11.8±0.8 Valencic & Smith 2015 R 7.6 Coe et al. 1996 A 7.6±1.0

16±1.8-

36.0±3.0

Goossens et al. 2013 R

26.0±2.0 Sguera et al. 2007 R

IGR

J2000.6+3210

HMXB 6.0b Marshall et al. 2006 13.6±2.0 Pradhan et al. 2013 R ∼4 Masetti et al. 2008 A 4.0±0.5

Sale et al. 2014 22±2.0 Morris et al. 2009 R

IGR

J21343+4738

HMXB 10.8±1.5 Reig & Zezas 2014 R 2.4±0.1 Reig & Zezas 2014 A 2.4±0.1

IGR

J22534+6243

HMXB 18.1±1.9 Esposito et al. 2013 A 18.1±1.9 5.6±0.3 Esposito et al. 2013 A 5.7±0.4

6.7 Masetti et al. 2013 A

KS 1947+300 HMXB 10.0 Liu et al. 2006 5.1±0.2 Epili et al. 2016 A 5.1±0.2 3.4 Negueruela et al. 2003 A 3.4±0.4

LS 2883 HMXB 2.7b Marshall et al. 2006 5.8±0.6 Tam et al. 2015 A 4.0±0.3 3.25 Westerlund & Garnier 1989 A 3.3±0.4

4.2±0.3 Pavlov et al. 2015 A

3.0±0.3 Kargaltsev et al. 2014 A

2.5±0.6 Pavlov et al. 2011 A

3.2±0.3 Chernyakova et al. 2006b A

5.0±0.3 Hirayama et al. 1999 A

5.6±0.6 Hirayama et al. 1996 A

LS 5039 HMXB 2.9±0.3 Ribo et al. 2002 6.4±0.6 Durant et al. 2011 A 6.4±0.1 3.8±0.2 Ribo et al. 2002 A 3.7±0.3

7.1±0.4 Rea et al. 2011 A 3.7±0.3 Clark et al. 2001 A

6.8±0.2 Zabalza et al. 2011 A 3.2±0.7 Motch et al. 1997 A

6.3±0.1 Bosch-Ramon et al. 2007 A

10±4 Reig et al. 2003 A

LS 992 HMXB 5±3 Reig & Roche 1999 R 2.0±0.2 Reig & Roche 1999 A 2.1±0.2

2.5±0.6 Motch et al. 1997 A

ALS 19596 HMXB 23.0±1.3 Piraino et al. 2000 A 23.0±1.3 8.1±0.9 Israel et al. 2001 A 8.1±0.9

MXB 1746-20 LMXB 8.5±0.4 Liu et al. 2006 6.5±0.2 Valencic & Smith 2015 A 6.7±0.3 A 5.7±0.1 3.3±0.3 Harris 2010 and 1996 A 3.2±0.3

7.4±0.4 Walsh et al. 2015 A 3.2±0.3 Minniti 1995 A

c©0000

RA

S,M

NR

AS

000,

000–000

20

Zh

u,

Tia

n,L

i&

Zh

an

g

SAX

J1748.9-2021

HMXB 8.5±0.4 Liu et al. 2006 7±1 Cackett et al. 2005 A 7±1 5.7±0.1 Pintore et al. 2016 3.1±0.3 Ortolani et al. 1994a A 3.2±0.4

3.3 Armandroff 1989 A

3.6 Harris & van den Bergh 1974 A

RX

J0146.9+6121

HMXB 2.3 Liu et al. 2006 4.9±0.5 La Palombara & Mereghetti 2006 A 4.9±0.5 2.9±0.1 Reig et al. 1997 A 2.9±0.1

9.4±1.2 Haberl et al. 1998a R 3.4 Motch et al. 1997 A

7.8±1.9 Haberl et al. 1998b R 2.9 Coe et al. 1993 A

RX

J0440.9+4431

HMXB 3.3 Liu et al. 2006 3.3±0.1 Ferrigno et al. 2013 A 3.4±0.1 2.0±0.2 Reig et al. 2005b A 2.2±0.2

7.3±1.5 La Palombara et al. 2012 A 2.8 Motch et al. 1997 A

4.7±0.4 Tsygankov et al. 2012 A

3.1±0.4 Usui et al. 2012 A

RX

J1829.4-2347

LMXB 6 Liu et al. 2006 2.2±0.1 Filippova et al. 2014 A 2.6±0.1 1.2 Barret et al. 1995 A 1.2±0.4

1.6±0.5 Thompson et al. 2005 A

4.0±0.4 Kong et al. 2000b A

1.1±0.2 in’t Zand et al. 1999 A

3.4±0.1 Thompson et al. 2008 A

RX J1832-33 LMXB 9.3 Liu et al. 2006 0.46±0.18 Parmar et al. 2001 A 0.55±0.22 0.58±0.06 Sidoli et al. 2008 A 0.58±0.06 0.3 Harris 2010 and 1996 A 0.3±0.1

0.89±0.35 Mukai & Smale 2000 A 0.3±0.1 Ortolani et al. 1994b A

Swift

J1753.5-0127

LMXB ∼6 Liu et al. 2007 2.5±0.1 Valencic & Smith 2015 A 2.2±0.1 2.8 Yoshikawa et al. 2015 A 2.7±0.1 1.4±0.1 Rahoui et al. 2015 A 1.3±0.1

2.7±0.1 Froning et al. 2014 A 4.1±1.4 Rushton et al. 2016 R 1.4±0.2 Froning et al. 2014 A

2.3±0.1 Reynolds et al. 2010 A 2.6±0.1 Tomsick et al. 2015 A 1.3±0.1 Durant et al. 2009 A

1.5±0.1 Reis et al. 2009 A 2.8±0.1 Mostafa et al. 2013 A 1.1±0.1 Cadolle Bel et al. 2007 A

2.3±0.2 Durant et al. 2009 A

1.7±0.1 Hiemstra et al. 2009 A

2.3±0.1 Miller et al. 2006 A

V404 Cyg LMXB 3.5 Liu et al. 2007 7.4±0.7 Bernardini & Cackett 2014 A 8.1±0.3 8.6±0.6 Rana et al. 2016 A 8.6±0.6 4.0 Hynes et al. 2009 A 3.6±0.5

6.1±0.6 Kong et al. 2002 A 3.1±0.9 Shahbaz et al. 2003 A

8.4±0.2 Plotkin et al. 2017 A 4.0 Casares et al. 1993 A

6.1±0.7 Hynes et al. 2009 A 3.4±0.6 Gotthelf et al. 1992 A

8.8±0.6 Bradley et al. 2007 A 3.2 Wagner et al. 1991 A

XTE J1550-564 LMXB 5.3 Liu et al. 2007 8.8±1.1 Corbel et al. 2006 A 6.6±0.1 6.6±0.1 Steiner et al. 2011 A 6.6±0.1 3.2±0.2 Curran & Chaty 2013 A 2.9±0.2

6.8±0.2 Kubota & Done 2004 A 2.2±0.3 Sanchez-Fernandez et al. 1999 A

6.5±0.1 Gierlinski & Done 2003 A

8.0±0.4 Miller et al. 2003 A

8.5±2.3 Tomsick et al. 2001 A

XTE J1720-318 LMXB 17.8±0.2 Valencic & Smith 2015 A 15.1±0.2 7±1 Chaty & Bessolaz 2006 A 6.6±0.6

12.4±0.2 Cadolle Bel et al. 2004 A 6.5±0.5 Nagata et al. 2003 A

XTE J1808-369 LMXB 3 Liu et al. 2007 2.1±0.1 Papitto et al. 2009 A 1.4±0.1 1.6±0.2 Patruno et al. 2009 A 1.4±0.1 0.8±0.3 Elebert et al. 2009 A 0.8±0.3

1.2±0.1 Heinke et al. 2009 A 1.3±0.1 Campana et al. 2008 A 0.7±0.3 Wang et al. 2001 A

0.9±0.1 Cackett et al. 2009b

1.5±0.4 Heinke et al. 2007

XTE?J1819-254 HMXB 9.5 Liu et al. 2006 2.4±0.4 Pahari et al. 2015 A 2.4±0.4 1.1±0.6 MacDonald et al. 2014 A 0.8±0.1

0.7±0.2

or

1.6±0.3

Lindstrom et al. 2005 A

0.8 Chaty et al. 2003 A

1.0±0.3 Orosz et al. 2001 A

0.7 Wagner 1999 A

XTE J1859+226 LMXB 7.6 Liu et al. 2007 ∼3 Wood et al. 1999 A 3.6±0.4 1.8±0.4 Hynes et al. 2002 A 1.8±0.4

∼8 dal Fiume et al. 1999 A

XTE J2103+457 HMXB 6.5 Liu et al. 2006 7.6±0.1 Camero et al. 2014 A 7.6±0.1 4.2±0.3 Hynes et al. 2004 A 4.2±0.3

7.0±0.3 Inam et al. 2004 A

37.6±1.2 Baykal et al. 2002 R

XTE J2123-058 LMXB 8.5 Liu et al. 2007 0.7+4.0−0.7

-

0.8+3.8−0.8

Tomsick et al. 2004 A 0.8±2.3 0.37±0.15 Hynes et al. 1998 A 0.37±0.15

>7.9 Tomsick et al. 1999 R

a : We convert the measured Balmer decrement to extinction using AV ≈ 10.331

× log(

Hα3Hβ

)

× 0.664 × 3.1 (Mavromatakis et al. 2004a).

b : The distances are estimated based on AV and the extinction maps of Marshall et al. (2006), Sale et al. (2014) and Green et al. (2014).

c©0000

RA

S,M

NR

AS

000,

000–000


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Documents

The gas-to-extinction ratio and the gas distribution in ...We investigate the relation between the optical extinction (AV) and the hydrogen column density (N H) determined from X-ray