EDP Sciences
Free Access
Volume 579, July 2015
Article Number A91
Number of page(s) 35
Section Galactic structure, stellar clusters and populations
DOI https://doi.org/10.1051/0004-6361/201424802
Published online 03 July 2015

Online material

Appendix A: Grouping of ATLASGAL sources

We give in this section details about our analysis to group sources, which are coherent in space and velocity using the friends-of-friends algorithm (Huchra & Geller 1982; Moore et al. 1993; Berlind et al. 2006). As a first step we sort ATLASGAL sources with known velocities (see Sect. 2.2) in the first and fourth quadrant according to an increasing galactic longitude range. For each ATLASGAL source we search for related sources in the whole sample according to a predetermined threshold in position and velocity. Associated ATLASGAL sources are added to the same group until no more sources that lie within a maximum distance and velocity interval to the neighbouring sources are found. The group is then considered to be complete and a new search around the next ATLASGAL source of the catalogue begins. All distances between at least two individual sources within one group must not exceed the maximum distance and the velocities of these must be within a given velocity interval. However, the size and the maximum difference of velocities within a group can exceed the input values for the maximum distance and the velocity interval used to search for associated sources. This algorithm results in an arrangement of the ATLASGAL sample into groups and sources that do not belong to any group.

We study the dependency of the grouping on the input parameters by varying the maximum distance and the velocity interval. This results in various divisions of the ATLASGAL sources into groups, which are analysed to improve the matching between our grouping and properties of known molecular cloud complexes (see Appendix B.2). For different input values for the maximum distance, Δdi, and a fixed interval of 10 km s-1 for velocities of neighbouring sources in a group we plot histograms of the size of a group in Fig. A.1. The size is defined as l × Δb)1 / 2 with the maximum difference in longitude, Δl, and in latitude, Δb, of sources within a group. For Δdi = 0.05° or 0.1° the size of most groups is smaller than those values, while the size of a larger number of groups is in the range of the input value Δdi = 0.2° or 0.3°. If Δdi is larger, 0.5°, some groups will have sizes larger than Δdi. This trend is visible for the ATLASGAL samples in the first (upper panel of Fig. A.1) and fourth quadrant (lower panel of Fig. A.1), although it is more obvious for the groups in the fourth quadrant. To avoid gathering separate groups together, we do not take a large input value such as 0.5°, but choose 0.3°. In addition, the median size of GRS molecular clouds (Roman-Duval et al. 2009) of ~0.3° is also consistent with the maximum distance of ATLASGAL sources within a group, which corresponds to ~22 pc at a median distance of ~4 kpc. Using 12CO emission García et al. (2014) obtained even sizes up to 100 pc for giant molecular clouds in the fourth quadrant.

thumbnail Fig. A.1

Relative number distribution of ATLASGAL groups with the size of a group for different input values of the maximum distance between individual sources within a group, Δdi. This is varied between 0.05° and 0.5° with a fixed interval for velocities of neighbouring sources in a group, Δvi, of 10 km s-1 for the sample in the first quadrant in panel a) and the groups in the fourth quadrant in panel b).

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thumbnail Fig. A.2

Relative number distribution of ATLASGAL groups in the first quadrant in panel a) and groups in the fourth quadrant in panel b) with the maximum velocity difference in a group for different input values of the velocity interval to search for associated sources in a group, Δvi. We use 5 to 15 km s-1 for that and a fixed maximum distance between individual sources in a group, Δdi, of 0.3°.

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Figure A.2 illustrates histograms of the maximum difference in velocities within a group for varying input values of the velocity interval used to search for neighbouring sources within a group, Δvi, and a fixed maximum distance between individual sources in a group. Figure A.2 indicates that with an increase in Δvi from 10 to 15 km s-1 the maximum velocity difference in a few groups becomes greater than the chosen Δvi. This is similar to the trend seen in the maximum distance, but less obvious. Since the maximum velocity difference of most groups should not exceed the input value of Δvi, we take 10 km s-1 for that. This analysis lets us choose the arrangement of ATLASGAL sources into groups resulting from 0.3° and 10 km s-1 as input values for the maximum distance and the maximum difference in velocities of neighbouring sources in a group. The maximum velocity difference of ATLASGAL sources in a group also agrees with the velocity dispersion of GRS molecular clouds (Roman-Duval et al. 2009) ranging up to 10 km s-1 and with the median linewidth of GMCs in the fourth quadrant (García et al. 2014) of 9 km s-1.

We call sources that are associated by their position and velocity “a group” and refer to the group’s spatial structure in a comparison with the ATLASGAL dust continuum emission as “a complex”. We illustrate an example of a complex within 22.1°<l< 22.6° and 0.25°<b< 0.5°, which results from our method, in the top panel of Fig. 2.

This method results in 296 groups in the first quadrant and 393 groups in the fourth quadrant, an uncertainty associated with these identifications might result from gathering separated groups along the line of sight together into one group.

thumbnail Fig. A.3

Number distribution of ATLASGAL groups with the number of sources per group in a), the velocity dispersion in b), and radius in c).

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Appendix A.1: Properties of groups

To characterize the groups we study some physical properties such as the velocity dispersion and physical radius. The distribution of the number of sources per group is illustrated in panel a) of Fig. A.3, the histogram of the velocity dispersion in panel b) and of the physical radius of the groups in panel c) using distances determined in Sect. 5. We obtain ~300 sources in the first and fourth quadrants that within 0.3° and 10 km s-1 are not related to another ATLASGAL source, as indicated in panel a) of Fig. A.3. The histograms of the velocity dispersion and physical radius of the groups containing 2 sources show that most groups exhibit a velocity dispersion of <2 km s-1 and radius of ~5 pc. The peak in the radius distribution mainly results from a subsample of about 100 groups, which contain only two sources per group. In contrast, the peak of the velocity dispersion is a property of the whole sample. Figure A.3 reveals a large dispersion of radii, which range to 80 pc. We obtain 12 groups out of the whole sample of 689 with radii larger than 35 pc, most of them consist of more than ten sources. Some of the large radii belong to groups at the far distances.

The correlation of the physical radius and velocity dispersion of all 709 groups is shown in Fig. A.4. A power-law relationship between the two properties is indicated as black line, which is given by Larson (1981)(A.1)where δυ is the velocity dispersion and R the radius. Some of the groups are located close to the relationship, but many groups lie in a broad distribution. To describe the whole sample we added a relation for a gravitationally bound cloud, resulting from a re-examination of Larson’s law by Ballesteros-Paredes et al. (2011)(A.2)with the H2 column density Σ. We plot Eq. (A.2) for H2 column densities of 1021 cm-2 as green line, 1022 cm-2 as red line, and 5 × 1022 cm-2 as blue line in Fig. A.4. The distribution of radius and velocity dispersion can be fitted by a power law with an exponent of 0.4 ± 0.02, which is shown as purple line. The slope is thus consistent with the power-law exponent given in Larson (1981). We performed a Spearman correlation test for an ATLASGAL subsample with heliocentric distances between 2 and 5 kpc, which gives a correlation coefficient of 0.57 with a p-value <0.0013. This indicates that the correlation between radius and velocity dispersion is significant over 3σ. A few groups are below the green line and exhibit extremely small velocity dispersions, <0.1 km s-1, because they mostly include only fewer than five sources with observed velocities. Because these do not have good statistics, we aim at describing groups associated with a larger number of measured velocities. A contour plot is presented in the upper panel of Fig. A.4 with a binning of 0.3 for the logarithm of the radius and for the logarithm of the velocity dispersion. It shows a trend of growing radii with increasing velocity dispersions and most groups follow the law of Ballesteros-Paredes et al. (2011).

thumbnail Fig. A.4

Correlation plot of the radius and velocity dispersion of all groups. The black line indicates the relation by Larson (1981), the green, red, and blue lines show the formula by Ballesteros-Paredes et al. (2011) for an H2 column density of 1021 cm-2, 1022 cm-2, and 5 × 1022 cm-2. The purple line shows a power-law fit to the data with an exponent of 0.4 ± 0.02. The contour plot is shown in the upper panel with the range of the logarithm of the radius and the velocity dispersion divided into bins of 0.3.

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Appendix B: Notes on individual sources

Appendix B.1: Sources with extreme properties

Clumps with high masses

The comparison of the kinematic distance with the gas mass in Fig. 8 reveals some extreme sources clustering at different distances that are known regions of massive star formation. Two clumps exhibiting a mass of 3.6 × 104M at a distance of 4.5 kpc are in the high mass star forming complex W33 (Bieging et al. 1978). Our KDA resolution yielding the near distance is consistent with trigonometric parallax measurements resulting in a distance of 2.53 kpc (Immer et al. 2013). The complex of sources at 5.5 kpc with high masses between 2 × 104 and 5 × 104M is located in W51 main. A similar distance of 5.4 kpc is derived from the trigonometric parallax (Sato et al. 2012) (see Appendix B.2). A cluster of ATLASGAL clumps at a distance around 8 kpc with gas masses higher than 104M are in the W43 complex and W43 south (see Appendix B.2). In addition, the source with a mass of 5.6 × 104M at 11 kpc is G10.47+0.03, an ultracompact HII region (Wood & Churchwell 1989) in the W31 complex (see Appendix B.2).

Sources with large distances from the Galactic plane

To investigate if individual complexes dominate extreme values for the distance from the Galactic mid-plane, an illustration of its radial distribution is shown for the ATLASGAL sample in the first quadrant in the top panel and for sources in the fourth quadrant in the lower panel of Fig. 15. Among the sources, which exhibit extremely high distances from the mid-plane indicated as red triangles, are the clumps at RGal ≈ 5.4 kpc and exhibiting a height larger than 200 pc below the plane in the upper panel. These are located in complex 4 at l ≈ 6.02° and b ≈ −1.23° (see Table 3). Because none of these sources contains an embedded HII region, we cannot analyse the HI absorption and determine the far kinematic distance from missing HI self-absorption, which is more uncertain than using the two methods. The complex is also known as NGC 6530, which is a young open cluster, for which Loktin & Beshenov (2001) derived a distance between 560 and 711 pc from Hipparcos trigonometric parallaxes, which would locate the region at a smaller height below the disk. We can also use only the HISA method for the clump at RGal ≈ 6.8 kpc with a height of 154 pc below the plane, located in complex 262 at l = 49.044° and b ≈ −1.078°, for which we determine the far kinematic distance. This source is associated with IRAS 19230+1341 and 6.7 −GHz methanol maser emission (Pestalozzi et al. 2005). Moreover, there is a complex of sources at RGal ≈ 5 kpc, which are located at a distance of 158 pc below the mid-plane. One of them at l = 17.608° and b = −0.726°, also known as IRAS 182271358, is in complex 75, which consists only of this source. The other two clumps at l ≈ 18.30° and b ≈ −0.72° are in complex 80, which contains no embedded HII region. We thus locate them at the far distance from missing HI self-absorption. Because the displacement from the Galactic plane of star forming regions is confined to 120 pc (Urquhart et al. 2011), the assignement of the far distance is unrealistic and the complexes are likely at the near distance. In addition, the HISA method is less reliable for the far distance than for the near distance. Moreover, the sources at a large height above the plane are not associated with 21 cm continuum sources either and we also use the HISA method to determine the far distance to them. Among those are the clump at RGal ≈ 4.6 kpc with a distance of ~185 pc from the disk, located in complex 58 at l ≈ 15.029° and b ≈ 0.851°, and one at RGal ≈ 5.5 kpc and a height of 178 pc, which is in complex 33 at l ≈ 11.92° and b ≈ 0.753°. Approximately 7% of the sources in the first and fourth quadrant, which we assign to the far distance, have a displacement larger than 120 pc from the plane. We keep the far distance for these clumps to use a consistent method to resolve the KDA for all ATLASGAL sources. Statistically this small fraction does not influence the Galactic distribution.

Clumps in the fourth quadrant (see lower panel of Fig. 15) with an extremely high displacement from the plane of 210 pc and a galactocentric radius of 5.8 kpc are in complex 187 (l = 330.044°,b = 1.050°). These sources are embedded in the IRDC G330.030+1.019 (Peretto & Fuller 2009), which hints at the near distance, while we determined the far distance for this complex. However, our KDA resolution is uncertain because we could use only the HISA method, which is less reliable for the far distance solution than for the near distance. Additional complexes, which we locate at the far distance based on this method and which are therefore at a large height above the Galactic plane are complex 186 (l = 329.718°,b = 0.804°) with only one source at RGal ≈ 5.8 kpc and a height of 160 pc. It is associated with the IRDC G329.717+0.785 (Peretto & Fuller 2009). Other examples are complex 280 (l = 340.088°,b = 0.928°) at RGal ≈ 4.2 kpc and a distance from the plane of 180 pc, embedded in the IRDC G340.094+0.929 (Peretto & Fuller 2009), as well as complex 360 (l = 350.816°,b = 0.514°) at RGal ≈ 8.4 kpc and a displacement from the plane of 150 pc, located within the IRDC G350.816+0.513 (Peretto & Fuller 2009).

Complexes with extremely large scale heights

The ATLASGAL sample within a galactocentric radius between 4.5 and 6.5 kpc shows a flat scale height distribution in Fig. 17, resulting mainly from a few complexes, which contain several sources. Approximately 7% of the sources in the first quadrant in the galactocentric radius bins from 4 to 7 kpc are in complex 261 (l = 49.2°,b = −0.7°) (see Appendix B.2) with a mean Galactic latitude of −0.31°. It is located at a kinematic distance of 5.54 kpc and at a mean height of 30.25 pc below the Galactic plane. Two additional large complexes in the first quadrant are close to the Galactic plane and thus have small heights: complex 172 (l = 30.767°,b = −0.050°) (see Appendix B.2) contains ~8% of the sources in the first quadrant with a mean latitude of 0.07°, a distance of 8.4 kpc, and a mean height of 10.67 pc. Approximately 8% of the clumps in the first quadrant are in complex 167 (l = 29.96°,b = −0.2°) (see Appendix B.2) with a mean latitude of −0.04°, a distance of 8 kpc, and a mean height of 5.84 pc below the mid-plane. In the fourth quadrant complex 242 (l = 336.97,b = −0.010) consists of ~10% of the sources within a galactocentric radius between 4 and 7 kpc. It is located at a mean latitude of 0.01°, nearby at a kinematic distance of 4.7 kpc, and thus has a small mean height of 1 pc below the Galactic plane. Complex 210 (l = 333.19°,b = −0.36°) (see Appendix B.2) has the largest contribution, ~12%, to the clumps in the fourth quadrant in the bins from 4 to 7 kpc. It exhibits a mean latitude of −0.38° and a distance of 3.52 kpc, which results in a mean height of 23.37 pc below the mid-plane.

Sources with high surface densities

The effective radius is compared with the gas mass of the ATLASGAL sample in Fig. 23. Some sources are located close to the dotted blue line, which indicates a surface density of ~1 g cm-2. Among those sources with high gas masses concentrated to small source radii is G35.20.74 with a mass of 423 M and an effective radius of about 0.1 pc, which is associated with compact HII regions detected at 6 cm by Dent et al. (1984). Another clump at l = 348.604°,b = −0.912°, which has a mass of 114 M and a radius of 0.08 pc, also exceeds the surface density of 1 g cm-2. It is an UCHIIR, which is also observed in CS (21) by Bronfman et al. (1996). One source with a higher mass of ~5 × 104M and a radius of ~1.5 pc is located in the W51 giant molecular cloud complex, associated with W51 IRS1 in W51 Main (see Appendix B.2). It emits at 20 μm, 2 cm, and 6 cm, which shows that it is an extended HII region illuminated by early OB stars in the centre (Genzel et al. 1982). Another clump with a slightly smaller mass of ~3 × 104M and effective radius of 1.24 pc is in W51 IRS2, which is a compact infrared source near W51 IRS1 (Genzel et al. 1982).

Appendix B.2: Identification of complexes with known giant molecular clouds

W43 Main

Complex 172 has most ATLASGAL sources in the first quadrant and consists of 47 clumps . It is located in the well known W43 main star forming complex (l = 30.767°,b = −0.050°, Motte et al. 2003), which contains a giant HII region. The association of giant molecular clouds with the W43 stellar cluster and the star formation efficiency of the molecular cloud complex of ~25% (Motte et al. 2003), which might result in a stellar density of ~100 stars pc-3, hint at a sudden ministarburst in W43. It can therefore be classified as a Galactic mini-starburst region and serve as a template to constrain cloud and gas properties in distant starburst galaxies (Motte et al. 2003).


We place 30 sources in the first quadrant in complex 167 (l = 29.96°,b = −0.2°, Nguyen Luong et al. 2011), which is W43-south. It is located at a slightly smaller Galactic longitude, 29.8°<l< 30.5°, than the W43 main cloud (complex 172). Investigation of 13CO emission toward the two (Nguyen Luong et al. 2011) revealed that they might be connected by low-density gas and that they likely belong to the same giant molecular cloud. The distance to the W43 region has already been derived by several studies: Pratap et al. (1999) obtained the near distance to the HII region G29.960.02, which is located in W43-south, by analysing extinction as well as formaldehyde absorption lines. Moreover, Motte et al. (in prep.) also derive the near distance to that source from large-scale HI absorption. Anderson & Bania (2009) determine the near distance to the W43 main cloud (complex 172) from HI self-absorption, while they locate G29.960.02 and another source in W43-south at the far distance from their investigation of HI self-absorption as well as HI absorption. These methods are also used by Russeil et al. (2011), who determine the near distance to several sources in the W43 main cloud and note that there is a disagreement in distances to G29.960.02 given in the literature. They use the near distance to that source derived from the maser parallax method. The W43 complex has also been observed by the BeSSeL Survey (Reid et al. 2009a): from observations of three 12 GHz methanol masers and a 22 GHz water maser Zhang et al. (2014) derived a parallax distance of 5.5 ± 0.37 kpc. Kolpak et al. (2003) obtained the far distance to two sources in W43-south by analysing HI absorption. Using the same method Pandian et al. (2008) derived the far distance to some clumps located in the W43 main cloud. These distance assignments agree with our distance resolutions for the W43 main cloud (complex 172) and W43-south (complex 167), which we place at the far distance. Investigation of HI self-absorption (see Sect. 5.2) toward the two complexes hints at the near distance, while HI absorption (see Sect. 5.4) toward four sources in the W43 main cloud and two clumps in W43-south reveals the far distance. Combining the two methods as described in Sect. 5.5 gives the far distance of 8.4 ± 0.82 kpc to complex 172 and 7.9 ± 0.64 kpc to complex 167.


We locate 22 sources in complex 50, which is M17 (the Omega Nebula, the Swan Nebula) (l = 15.10°,b = −0.75°Wu et al. 2009). It is one of the giant molecular clouds with the largest sizes and masses in the Milky Way (Lada 1976). A study of dust emission at near-infrared, far-infrared, and submm wavelengths (Chini et al. 1987) showed that M17 is associated with one of the most luminous compact HII regions in the Galaxy. Lada & Chaisson (1975) observed 6-cm H2CO aborption lines toward M17, which revealed a lane of dust obscuration west of the HII region. This part of the molecular cloud contains a small region of strong CO emission located south-west of the continuum peak at 6 cm (M17 SW) and is also detected in other molecular lines such as HCN and H2CO at 140 GHz (Lada et al. 1974) probing high cloud densities of 105 cm-3 (Thaddeus et al. 1971). Varying H2CO linewidths (Lada & Chaisson 1975) and CO measurements in M17 SW, which give a lower mass limit of 6000 M (Lada et al. 1974), hint at a molecular cloud collapse. The 12 GHz CH3OH maser line emitted in M17 SW was measured by Xu et al. (2011) to determine the trigonometric parallax. They obtain a distance of 2 ± 0.13 kpc and locate G15.030.68 in the Sagittarius spiral arm. We also derive a near distance of 2.4 ± 0.61 kpc from HI self-absorption and HI absorption toward two sources in the complex.

W51 Main

Complex 261 with 22 ATLASGAL clumps is part of the W51 molecular cloud (l = 49.2°,b = −0.7°, Green 2009), known as W51 Main. The whole W51 complex is one of the most massive (1.2 × 106M) and largest (Δl × Δb = 83 × 114 pc) giant molecular clouds (Carpenter & Sanders 1998). In addition, W51 is one of the most luminous (107L) star forming complexes and harbours many O stars (Bieging 1975), which indicate that the region is currently forming an OB association. Moreover, Kumar et al. (2004) discovered embedded star clusters from their near-infrared observations, which are associated with UCHIIRs. They obtain a high star formation efficiency of ~10% from an estimate of the W51 molecular cloud mass. Sato et al. (2012) observed the 22 GHz H2O maser line in W51 Main to measure the trigonometric parallax, which results in a distance of ~5.41 kpc. The 21 cm continuum data toward four ATLASGAL sources show absorption lines at velocities between 65 and 73 km s-1, which are consistent with the higher velocity range of some of the absorption spectra in the W51 molecular cloud by Koo (1997). Previous studies placed the complex in the Sagittarius arm at the tangent point (Koo 1997; Etoka et al. 2012), which agrees with our distance assignment at 5.5 ± 1.67 kpc.


We place 21 sources in complex 22, which is W31 (l = 10.3262°,b = −0.1432°, Goss & Shaver 1970). It is a massive star forming region, which was resolved into two components (G10.20.3 and G10.30.1) by Goss & Shaver (1970) and Shaver & Goss (1970) using observations at 5 GHz and 408 MHz. Complex 22 also includes two UCHIIRs (G10.150.34 and G10.300.15, Wood & Churchwell 1989) located in the larger scale surrounding of the two subregions. Downes et al. (1980) compared velocities of the H110α hydrogen recombination lines of clumps in W31 with H2CO absorption line velocities to resolve the KDA, which results in the near kinematic distance of 6 kpc. Blum et al. (2001) measure a spectrophotometric distance of 3.4 kpc, while Wilson (1972) obtains H2CO absorption lines at velocities, which are higher than the hydrogen recombination line velocity, and therefore assigns the complex to the far distance. Georgelin & Georgelin (1976) also measure absorption at velocities exceeding the velocity of the HII regions, which hints at the far distance of 18.7 kpc. In addition, analysis of H2CO, OH, and HI absorption lines by Corbel et al. (1997) indicates the far kinematic distance of 14.5 kpc. Examination of HI self-absorption toward the complex yields the near distance. In addition, HI absorption toward two clumps (G10.150.34 and G10.320.16) is at velocities, which are larger than the source velocity, but much smaller than the tangent point velocity and therefore hints at the near distance of 2 ± 0.93 kpc. However, this analysis cannot reject the far distance because absorption at velocities larger than ~50 km s-1 is unlikely given the lack of HI and CO gas at Galactic longitudes between 5° and 25° (see e.g. Fig. 3 in Dame et al. 1987; Corbel & Eikenberry 2004). Sanna et al. (2014) observed 22.2 GHz H2O maser emission toward G10.6200.38 to measure the trigonometric parallax within the BeSSeL Survey (Reid et al. 2009a), which resulted indeed in a distance of 5 kpc and located W31 on the near edge of the hole in the gas distribution within a galactocentric radius of ~3 kpc.


Complex 156 consisting of 17 ATLASGAL clumps is associated with the HII region N49 (l = 28.827°,b = −0.229°, Churchwell et al. 2006). It is one of the bubbles, which are detected by the investigation of GLIMPSE (Benjamin et al. 2003) images from Churchwell et al. (2006). This region consists of a layer of dense neutral material, which accumulates as a shell around a radio HII region with an approximately spherical shape. 20 cm free-free emission probes the ionized gas of N49, the thermal continuum at 24 μm traces the hot dust within the HII region, which is surrounded by an outer layer of 8 μm emission (Watson et al. 2008). A hole is revealed in the centre of the 24 μm and 20 cm emission, where the exciting star is located. This might be an O5V or O8III star, although the age estimate by Watson et al. (2008) favours the O5V star with 105 yr. Analysis of the 870 μm emission yields four condensations within the shell around the central HII region and a mass of 4200 M for the whole shell (Deharveng et al. 2010). NH3 velocities (Wienen et al. 2012) are observed toward three condensations and the 13CO velocity (Jackson et al. 2006) is measured toward the fourth condensation indicating their assocation with the ionized gas. Two massive young stellar objects can be revealed at 24 μm, which are in an early evolutionary phase (Watson et al. 2008), as well as an UCHIIR located in the surrounding of the most massive condensation (Deharveng et al. 2010). Anderson & Bania (2009) derived the near distance to the region from HI self-absorption and HI absorption. Because none of the observed ATLASGAL sources in N49 satisfied our criteria for radio continuum absorption, we could only look for HI self-absorption and we also determined the near distance of 5.7 ± 0.41 kpc.


We locate ten sources in complex 70, which is M16 (the Eagle Nebula) (l = 16.94°,b = 0.76°, Zeilik & Lada 1974). The HII region M16 belongs to a larger star forming complex, which is mostly still embedded within the molecular cloud (Goudis 1976). Images taken by the Hubble Space Telescope reveal structures in the molecular cloud containing dense globules, which are photoevaporated more slowly than their lower density environment, forming “molecular pillars”. The heating and ionization of M16 results from the young open star cluster NGC 6611. The Eagle Nebula is associated with recent and ongoing high mass star formation. Its stellar population has been analysed by Hillenbrand et al. (1993). They estimated an age of (2±1) ×106 yr and locate the region at a distance of 2 kpc from spectroscopic parallaxes. To resolve the KDA we could only analyse HISA because complex 70 contains no 21 cm continuum source. We obtain a near distance of 2.2 ± 0.56 kpc, which agrees with the distance derived by Hillenbrand et al. (1993).


Complex 210 consists of most ATLASGAL sources in the fourth quadrant containing 58 clumps (l = 333.19°,b = −0.36°, Bains et al. 2006). It is known as the G333 giant molecular cloud complex or RCW106 region with several molecular clouds forming high mass stars as well as HII regions. We determine the near kinematic distance to complex 210 of 3.5 ± 0.35 kpc from HI self-absorption and HI absorption toward several sources. Our estimate is consistent with the kinematic distance of 3.6 kpc calculated by Lockman (1979) as well as the range from 3.3 to 3.8 kpc given by Urquhart et al. (2012) and with the spectrophotometric distance of 3.96 kpc obtained by Moisés et al. (2011).


From our ATLASGAL sample in the fourth quadrant 41 sources are located in complex 29 (l = 305.361°,b = 0.056°, Clark & Porter 2004), which is the rich high mass star forming complex G305. Some luminous HII regions are embedded within the molecular cloud complex, the analysis of MSX data by Clark & Porter (2004) yields strong radio sources with up to 31 O7V stars ionizing an HII region. Russeil (2003) and Moisés et al. (2011) divided the region into two parts, G305.2+0.0 and G305.2+0.2, and derived distances to them. While Russeil (2003) obtained a distance of 3.5 kpc for both, the spectrophotometric distance analysed by Moisés et al. (2011) is inconclusive for G305.2+0.0, but yields the same result as assigned by Russeil (2003) for G305.2+0.2. Our HI self-absorption method also indicates the near distance to complex G305. However, the analysis of HI absorption toward one source of the region, G305.19+0.03 (see Fig. B.1), results in the far distance. The composite of the two methods reveals the far distance of 6.6 ± 1.24 kpc to the complex.

NGC 6334

Complex 356 with 28 ATLASGAL clumps is known as the NGC 6334 molecular cloud complex (l = 351.317°,b = 0.661°, Rodriguez et al. 1982) with embedded HII regions identified at 6 cm as compact to extended radio sources. Neckel (1978) derived the distance to NGC 6334 from optical observations of young stars within this complex of 1.74 ± 0.31 kpc and place it in the Sagittarius arm. They also give a kinematic distance of 0.7 ± 1.9 kpc, which corresponds to the interarm region. The two values are within the uncertainty of the near kinematic distance of 0.8 ± 1.03 kpc, which we obtained. The large distance uncertainty originates from the unreliability of kinematic distances near a longitude of 0°. Our result also agrees with the spectrophotometric distance of 1.82 kpc given by Moisés et al. (2011) and with the distance of 1.35 kpc derived recently from trigonometric parallax measurement by Reid et al. (2014).


Complex 255 contains 24 ATLASGAL sources, which connects three small complexes. The complex is a giant HII region, where Caswell & Haynes (1987) measured hydrogen recombination lines near 5 GHz. Caswell & Haynes (1987) derived a vast range of recombination line velocities over 120 km s-1 indicating several features, which lie on the same line of sight to the observer at different distances. We obtain a narrower velocity range from −30 to −42 km s-1 of most sources and a few clumps exhibit −46 to −56 km s-1 in our smaller complex at a longitude of 338°. Georgelin et al. (1996) also determined velocities and distances to two radio sources in the complex from observations of Hα emission. They give kinematic distances of 2.2 kpc for 338.398+0.164 and 3.2 kpc for 338.450+0.061, which lie within the errors of our distance estimates. However, they should assign the far distance taking the OH absorption line into account. While Russeil (2003) also places the complex at the far distance, our KDA resolution reveals the near distance of 3.06 kpc. The HI intensity map shows absorption, which hints at the near kinematic distance. In addition, the HI absorption method results in the near kinematic distance to one HII region and the far distance to four 21 cm continuum sources. The combination of the two techniques gives the near kinematic distance of 3.1 ± 0.41 kpc to the complex.

Gum 50

We locate 23 sources in the fourth quadrant in complex 163, which is known as Gum 50 (l = 328.573°,b = −0.531°, Beichman et al. 1988). Observations of hydrogen recombination lines near 5 GHz by Caswell & Haynes (1987) reveal HII regions located in the complex such as IRAS 155395353 and G328.570.53. Moreover, a few sources in this high mass star forming region, e.g. IRAS 155255407, IRAS 155395353, and G328.340.53, are associated with IRDCs (Peretto & Fuller 2009), which indicates the near distance. This is in agreement with our HI self-absorption and HI absorption analysis resulting in the near kinematic distance of 3 ± 0.39 kpc.

Table B.1

Associations of largest ATLASGAL complexes with known giant molecular clouds (see Sect. 7.1).

thumbnail Fig. B.1

HI spectrum toward G305.19+0.03, an HII region in the high mass star forming complex G305 extracted from the SGPS. For explanation of the different lines see Fig. 3. The HI absorption method yields the far distance.

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Appendix C: Comparison of kinematic and trigonometric parallax distances

Recently, more than a hundred distances to high mass star forming regions have been determined via measurements of trigonometric parallaxes of class II CH3OH and H2O masers with Very Long Baseline Interferometry (VLBI). The data were taken with the VLBI Exploration of Radio Astronomy (VERA) array2 and with the NRAO Very Long Baseline Array in the course of Bar and Spiral Structure Legacy (BeSSeL) survey3 (Reid et al. 2014). Parallaxes provide gold standard distance measurements since they result from direct measurements and basic trigonometry and do not require assumptions such as a rotation curve.

Of the 103 sources for which Reid et al. (2014) present parallax distances, 48 are covered by ATLASGAL. For total of 27 of these we have determined kinematic distances. Our LSR velocities, determined from NH3 spectra agree with maser velocities within the uncertainties assigned by Reid et al. (2014), in 15 cases within 2 km s-1 and in all cases within 6 km s-1. In Fig. C.1 we compare the parallax distances with our kinematic distances.

For a third of the 27 sources the kinematic and the parallax distance agree within the errors and all but two within a factor of two. For the latter two, G12.88+0.48 and G16.580.05, choosing the far kinematic distances, 12.9 and 11.7 kpc, respectively, obviously was a wrong decision, which for the former is reinforced by its high latitude. Both are found in complexes that do not contain radio continuum sources and thus have no HI absorption. They were placed at the far distance because they also did not exhibit HI self-absorption.

Part of the scatter in Fig. C.1 may be caused by peculiar velocities, which of course are not accounted for in the rotation curve model used to determine the kinematic distances. A dramatic example for this is the case of W3OH for which the parallax distance is almost a factor of two smaller than the kinematic distance (Xu et al. 2006).

thumbnail Fig. C.1

Kinematic distances determined in our study vs. trigonometric parallax distances from the BeSSeL survey and VERA results. The dotted fiducial line marks equality.

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Some of the sources with measured parallax distance and resolved kinematic distance ambiguity are described in more detail.


G12.89+0.49 is in complex 44. This source is a high mass protostellar object (Beuther et al. 2002), which is associated with IRAS 180891732 (Sridharan et al. 2002). It is well known to harbour strong H2O (Anglada et al. 1996), CH3OH, (Caswell et al. 1995) and OH masers (Cohen et al. 1991). Previous studies (Sridharan et al. 2002; Beuther et al. 2002) give the near and far kinematic distance, but did not solve the KDA. Using 12.2 GHz methanol masers Xu et al. (2011) derived a distance of 2.34 kpc from the trigonometric parallax. Immer et al. (2013) also measured a parallax distance from 22.2 GHz water maser observations, which is in agreement with the result of Xu et al. (2011), and locate the high mass star forming region in the Scutum spiral arm. Because there is no 21 cm continuum source in complex 44, we could only analyse HISA, which reveals the far kinematic distance of 13 ± 0.46 kpc. However, the HISA method is less reliable for the far distance than for the near distance.

G12.910.26 and G12.680.18

These two sources are located in the W33 molecular cloud complex. It is a strong radio continuum source detected by Westerhout (1958) and consists of various evolutionary phases of high mass star formation from quiescent IRDCs to HII regions. Radio recombination line observations revealed two velocity components in different parts of W33, one at ~36 km s-1 and one at ~58 km s-1 (Bieging et al. 1978). We also obtain several velocities from NH3 observations in this region and divide it into different complexes. We locate G12.910.26 in complex 38 with a mean radial velocity of 43 km s-1 and a velocity dispersion of 6.81 km s-1 and G12.680.18 in complex 41 with a mean velocity of 55.56 km s-1 and a velocity dispersion of 0.57 km s-1. Immer et al. (2013) used a water maser exhibiting velocities between 36 to 38 km s-1 to derive the parallax of their observation at l = 12.909°,b = −0.261°, which corresponds to a distance of 2.53 kpc. This is consistent with our investigation of HI self-absorption and HI absorption towards two sources in complex 38, which reveal the near distance of 4.2 ± 0.39 kpc. Immer et al. (2013) also fit the position of a water maser to obtain the parallax of G12.680.18, which results in a distance of 2.4 kpc. This is in agreement with our KDA resolution of complex 41, which gives the near kinematic distance of 4.9 ± 0.3 kpc using only the HISA method.


G23.010.41 is located in complex 114. This high mass star forming region contains a 12 GHz methanol maser, which coincides with a H2CO maser and NH3 (3,3) peak emission. A fit to the positions of the two masers gives one parallax, which corresponds to a distance of 4.59 kpc (Brunthaler et al. 2009). Different distances to this source are given by previous studies so far: Codella et al. (1997), Forster & Caswell (1999) and Caswell & Haynes (1983) used the far kinematic distance, although the distance ambiguity was not resolved. In contrast, Harju et al. (1998) locate the region at the near kinematic distance. This is consistent with our KDA resolution resulting in 4.9 ± 0.29 kpc from HI self-absorption.


G23.460.18 belongs to complex 117. This is a star forming region, where 6.7 GHz methanol maser emission observed by Walsh et al. (1998) hints at massive young stellar objects. Previous studies revealed different velocity components of the molecular cloud harbouring G23.440.18: Ohishi et al. (2012) obtained three components at 101, 103, and 104 km s-1 from their H13CO+ spectra, while Lockman (1989) derived the radio recombination line velocity of 103 km s-1 for an HII region, which is likely associated with the molecular cloud. Rathborne et al. (2009) give two 13CO velocities at 101 km s-1 and ~104 km s-1, for which Roman-Duval et al. (2009) determined kinematic distances of 6.43 kpc to the component at 101 km s-1 and of 6.65 kpc to the 104 km s-1 component from HISA and HI absorption. Brunthaler et al. (2009) fitted the positions of four 12 GHz methanol masers to measure the parallax of their observation at l = 23.4398°,b = −0.1822°. This yields a distance of 5.88 kpc and places the region near the end of the galactic bar in the Norma arm. We measured a velocity of 98.4 km s-1 toward G23.460.18, which is close to the velocity of the methanol maser at 97.6 km s-1 (Brunthaler et al. 2009). We derive the near kinematic distance of 5.9 ± 0.3 kpc using the HISA method, which is in agreement with the distance from the trigonometric parallax and also similar to the kinematic distance of 6.43 kpc of Roman-Duval et al. (2009).


G23.960.11 is in complex 121, which is associated with G23.6570.127 observed with the BeSSeL survey. G23.6570.127 is a massive protostar, where Bartkiewicz et al. (2005) detected 6.7 GHz methanol maser emission between 77 and ~88 km s-1. Analysis of an EVN image showed that the methanol masers are arranged in an approximately circular ring (Bartkiewicz et al. 2005). Bartkiewicz et al. (2008) observed 12.2 GHz methanol masers with the VLBA towards G23.6570.127, which have all 6.7 GHz counterparts and are also distributed spherically symmetric. Fitting 19 positions of 12.2 GHz methanol masers, which have approximately the same velocities as the 6.7 GHz methanol masers, reveals a distance of 3.19 kpc from the trigonometric parallax (Bartkiewicz et al. 2008). We obtain a similar velocity range of complex 121 with a mean velocity of 78 km s-1 and a velocity dispersion of 3.7 km s-1. The HISA method and HI absorption towards one source in complex 121 results in the near kinematic distance of 4.9 ± 0.31 kpc, which is similar to the distance derived by Bartkiewicz et al. (2008).


G35.190.74 is located in complex 214 and associated with G35.20.7, which is observed by the BeSSeL survey. The complex is a massive star forming region, where Gibb et al. (2003) observed CO, SiO, and radio emission indicating several outflows from G35.20.7. In addition, VLA observations at 3.5 and 6 cm showed fragmentation of the radio continuum into multiple YSOs embedded within dense clumps along the outflow (Gibb et al. 2003). Solomon et al. (1987) derived the near distance of 3.3 kpc because the far distance would locate the source at a larger height than 150 pc from the plane. The positions of two 12 GHz masers with radial velocities of 27.9 and 27.5 km s-1 were fitted by Zhang et al. (2009) to obtain the trigonometric parallax, which leads to a distance of 2.19 kpc. The mean velocity of complex 214 is 34.1 km s-1 with a velocity dispersion of 2.02 km s-1, which agrees with the velocities of the masers. Because complex 214 contains no 21 cm continuum source, we use only the HISA method, which reveals the near kinematic distance of 2.3 ± 0.43 kpc. Our KDA resolution is therefore consistent with the distances given by Solomon et al. (1987) and Zhang et al. (2009).

© ESO, 2015

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