Physical and chemical conditions in methanol maser selected hot cores and UCH ii regions
Identifieur interne : 002862 ( Istex/Corpus ); précédent : 002861; suivant : 002863Physical and chemical conditions in methanol maser selected hot cores and UCH ii regions
Auteurs : C. R. Purcell ; S. N. Longmore ; M. G. Burton ; A. J. Walsh ; V. Minier ; M. R. Cunningham ; R. BalasubramanyamSource :
- Monthly Notices of the Royal Astronomical Society [ 0035-8711 ] ; 2009-03-21.
Abstract
We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67 GHz methanol maser selected star-forming clumps observed by Purcell. In addition to the previously reported measurements of HCO+ (1–0), H13CO+ (1–0) and CH3CN (5–4) and (6–5), we used the Mopra antenna to detect emission lines of N2H+ (1–0), HCN (1–0) and HNC (1–0) towards 82/83 clumps (99 per cent), and CH3OH (2–1) towards 78/83 clumps (94 per cent). The molecular line data have been used to derive virial and local thermodynamic equilibrium masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between subsamples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6 GHz radio emission, indicating the presence of an Ultra-Compact H ii (UCH ii) region, and ‘isolated’ masers (without associated radio emission), and between clumps exhibiting CH3CN emission and those without. In particular, thermal CH3OH is found to be brighter and more abundant in UCH ii regions and in sources with detected CH3CN, and may constitute a crude molecular clock in single dish observations. Clumps associated with 8.6 GHz radio emission tend to be more massive and more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Hyper-Compact H ii or UCH ii region is the first visible tracer of star formation. The gas mass to submm/infrared luminosity relation for the combined sample was found to be L∝M0.68, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than, the mass of the stellar content. We also find that the mass of the hot core is correlated with the mass of the clump in which it is embedded.
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DOI: 10.1111/j.1365-2966.2008.14283.x
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R." last="Purcell">C. R. Purcell</name><affiliation><mods:affiliation>University of Manchester, Jodrell Bank Observatory, Macclesfield, Cheshire SK11 9DL</mods:affiliation></affiliation><affiliation><mods:affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</mods:affiliation></affiliation><affiliation><mods:affiliation>E-mail: Cormac.Purcell@manchester.ac.uk</mods:affiliation></affiliation><affiliation><mods:affiliation></mods:affiliation></affiliation><affiliation><mods:affiliation>E-mail: Cormac.Purcell@manchester.ac.uk</mods:affiliation></affiliation></author><author><name sortKey="Longmore, S N" sort="Longmore, S N" uniqKey="Longmore S" first="S. N." last="Longmore">S. N. 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R. Cunningham</name><affiliation><mods:affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</mods:affiliation></affiliation></author><author><name sortKey="Balasubramanyam, R" sort="Balasubramanyam, R" uniqKey="Balasubramanyam R" first="R." last="Balasubramanyam">R. Balasubramanyam</name><affiliation><mods:affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</mods:affiliation></affiliation><affiliation><mods:affiliation>Raman Research Institute, Sadashivanagar, Bangalore 560 080, India</mods:affiliation></affiliation></author></analytic><monogr></monogr><series><title level="j">Monthly Notices of the Royal Astronomical Society</title><title level="j" type="abbrev">Monthly Notices of the Royal Astronomical Society</title><idno type="ISSN">0035-8711</idno><idno type="eISSN">1365-2966</idno><imprint><publisher>Blackwell Publishing Ltd</publisher><pubPlace>Oxford, UK</pubPlace><date type="published" when="2009-03-21">2009-03-21</date><biblScope unit="volume">394</biblScope><biblScope unit="issue">1</biblScope><biblScope unit="page" from="323">323</biblScope><biblScope unit="page" to="339">339</biblScope></imprint><idno type="ISSN">0035-8711</idno></series></biblStruct></sourceDesc><seriesStmt><idno type="ISSN">0035-8711</idno></seriesStmt></fileDesc><profileDesc><textClass></textClass></profileDesc></teiHeader><front><div type="abstract">We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67 GHz methanol maser selected star-forming clumps observed by Purcell. In addition to the previously reported measurements of HCO+ (1–0), H13CO+ (1–0) and CH3CN (5–4) and (6–5), we used the Mopra antenna to detect emission lines of N2H+ (1–0), HCN (1–0) and HNC (1–0) towards 82/83 clumps (99 per cent), and CH3OH (2–1) towards 78/83 clumps (94 per cent). The molecular line data have been used to derive virial and local thermodynamic equilibrium masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between subsamples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6 GHz radio emission, indicating the presence of an Ultra-Compact H ii (UCH ii) region, and ‘isolated’ masers (without associated radio emission), and between clumps exhibiting CH3CN emission and those without. In particular, thermal CH3OH is found to be brighter and more abundant in UCH ii regions and in sources with detected CH3CN, and may constitute a crude molecular clock in single dish observations. Clumps associated with 8.6 GHz radio emission tend to be more massive and more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Hyper-Compact H ii or UCH ii region is the first visible tracer of star formation. The gas mass to submm/infrared luminosity relation for the combined sample was found to be L∝M0.68, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than, the mass of the stellar content. We also find that the mass of the hot core is correlated with the mass of the clump in which it is embedded.</div></front></TEI><istex><corpusName>oup</corpusName><author><json:item><name>C. R. Purcell</name><affiliations><json:string>University of Manchester, Jodrell Bank Observatory, Macclesfield, Cheshire SK11 9DL</json:string><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string><json:string>E-mail: Cormac.Purcell@manchester.ac.uk</json:string><json:null></json:null><json:string>E-mail: Cormac.Purcell@manchester.ac.uk</json:string></affiliations></json:item><json:item><name>S. N. Longmore</name><affiliations><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string><json:string>Harvard–Smithsonian Center For Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA</json:string><json:string>CSIRO Australia Telescope National Facility, PO Box 76, Epping, NSW 1710, Australia</json:string></affiliations></json:item><json:item><name>M. G. Burton</name><affiliations><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string></affiliations></json:item><json:item><name>A. J. Walsh</name><affiliations><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string><json:string>Centre for Astronomy, James Cook University, Townsville, QLD 4811, Australia</json:string></affiliations></json:item><json:item><name>V. Minier</name><affiliations><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string><json:string>Service d'Astrophysique, DAPNIA/DSM/CEA Saclay, 91191 Gif-sur-Yvette, France</json:string><json:string>AIM, Unité Mixte de Recherche, CEA-CNRS-Université Paris VII, UMR 7158, CEA/Saclay, 91191 Gif-sur-Yvette, France</json:string></affiliations></json:item><json:item><name>M. R. Cunningham</name><affiliations><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string></affiliations></json:item><json:item><name>R. Balasubramanyam</name><affiliations><json:string>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</json:string><json:string>Raman Research Institute, Sadashivanagar, Bangalore 560 080, India</json:string></affiliations></json:item></author><subject><json:item><value>surveys</value></json:item><json:item><value>stars: formation</value></json:item><json:item><value>stars: pre-main-sequence</value></json:item><json:item><value>ISM: abundances</value></json:item><json:item><value>ISM: molecules</value></json:item></subject><arkIstex>ark:/67375/HXZ-7HBKJRZ7-G</arkIstex><language><json:string>unknown</json:string></language><originalGenre><json:string>research-article</json:string></originalGenre><abstract>We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67 GHz methanol maser selected star-forming clumps observed by Purcell. In addition to the previously reported measurements of HCO+ (1–0), H13CO+ (1–0) and CH3CN (5–4) and (6–5), we used the Mopra antenna to detect emission lines of N2H+ (1–0), HCN (1–0) and HNC (1–0) towards 82/83 clumps (99 per cent), and CH3OH (2–1) towards 78/83 clumps (94 per cent). The molecular line data have been used to derive virial and local thermodynamic equilibrium masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between subsamples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6 GHz radio emission, indicating the presence of an Ultra-Compact H ii (UCH ii) region, and ‘isolated’ masers (without associated radio emission), and between clumps exhibiting CH3CN emission and those without. In particular, thermal CH3OH is found to be brighter and more abundant in UCH ii regions and in sources with detected CH3CN, and may constitute a crude molecular clock in single dish observations. Clumps associated with 8.6 GHz radio emission tend to be more massive and more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Hyper-Compact H ii or UCH ii region is the first visible tracer of star formation. The gas mass to submm/infrared luminosity relation for the combined sample was found to be L∝M0.68, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than, the mass of the stellar content. We also find that the mass of the hot core is correlated with the mass of the clump in which it is embedded.</abstract><qualityIndicators><refBibsNative>true</refBibsNative><abstractWordCount>298</abstractWordCount><abstractCharCount>1862</abstractCharCount><keywordCount>5</keywordCount><score>10</score><pdfWordCount>12757</pdfWordCount><pdfCharCount>68680</pdfCharCount><pdfVersion>1.4</pdfVersion><pdfPageCount>17</pdfPageCount><pdfPageSize>595.274 x 782.286 pts</pdfPageSize></qualityIndicators><title>Physical and chemical conditions in methanol maser selected hot cores and UCH ii regions</title><genre><json:string>research-article</json:string></genre><host><title>Monthly Notices of the Royal Astronomical Society</title><language><json:string>unknown</json:string></language><issn><json:string>0035-8711</json:string></issn><eissn><json:string>1365-2966</json:string></eissn><publisherId><json:string>mnras</json:string></publisherId><volume>394</volume><issue>1</issue><pages><first>323</first><last>339</last></pages><genre><json:string>journal</json:string></genre><subject><json:item><value>Papers</value></json:item></subject></host><namedEntities><unitex><date><json:string>2009-03-07</json:string><json:string>2002</json:string></date><geogName></geogName><orgName><json:string>University of New South Wales School of Physics Scholarship</json:string><json:string>CSIRO Australia Telescope National Facility, PO Box</json:string><json:string>University of New South Wales</json:string><json:string>National Facility</json:string><json:string>Australian Research Council</json:string><json:string>Australia Telescope</json:string><json:string>James Cook University</json:string><json:string>Raman Research Institute, Sadashivanagar, Bangalore</json:string><json:string>Australia Telescope Compact Array</json:string><json:string>Harvard–Smithsonian Center For Astrophysics</json:string><json:string>University of Tasmania</json:string><json:string>Commonwealth of Australia</json:string></orgName><orgName_funder><json:string>Commonwealth of Australia</json:string></orgName_funder><orgName_provider></orgName_provider><persName><json:string>M. G. Burton</json:string><json:string>R. Balasubramanyam</json:string><json:string>The</json:string><json:string>Histograms</json:string><json:string>S. N. Longmore</json:string><json:string>M. R. Cunningham</json:string><json:string>C. R. Purcell</json:string><json:string>A. J. Walsh</json:string></persName><placeName><json:string>Australia</json:string><json:string>Townsville</json:string><json:string>India</json:string><json:string>Gif-sur-Yvette</json:string><json:string>MA</json:string><json:string>France</json:string><json:string>Cambridge</json:string></placeName><ref_url></ref_url><ref_bibl><json:string>Walsh et al. (1998)</json:string><json:string>Weidner & Kroupa 2004</json:string><json:string>Pickett et al. 1998</json:string><json:string>Cesaroni et al. 2005</json:string><json:string>Hill et al. (2005)</json:string><json:string>Goldsmith & Langer 1999</json:string><json:string>Millar, MacDonald & Gibb 1997</json:string><json:string>C. R. 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(1997)</json:string><json:string>Churchwell 2002</json:string><json:string>Cummins, Linke & Thaddeus 1986</json:string><json:string>Kim et al. 2000, 2001</json:string><json:string>Kurtz et al. 2000</json:string><json:string>Townes & Schawlow 1955</json:string><json:string>Solomon et al. 1987</json:string><json:string>Avalos et al. 2006</json:string></ref_bibl><bibl></bibl></unitex></namedEntities><ark><json:string>ark:/67375/HXZ-7HBKJRZ7-G</json:string></ark><categories><wos><json:string>science</json:string><json:string>astronomy & astrophysics</json:string></wos><scienceMetrix><json:string>natural sciences</json:string><json:string>physics & astronomy</json:string><json:string>astronomy & astrophysics</json:string></scienceMetrix><scopus><json:string>1 - Physical Sciences</json:string><json:string>2 - Earth and Planetary Sciences</json:string><json:string>3 - Space and Planetary Science</json:string><json:string>1 - Physical Sciences</json:string><json:string>2 - Physics and Astronomy</json:string><json:string>3 - Astronomy and Astrophysics</json:string></scopus><inist><json:string>sciences appliquees, technologies et medecines</json:string><json:string>sciences biologiques et medicales</json:string><json:string>sciences biologiques fondamentales et appliquees. psychologie</json:string><json:string>biophysique des tissus, organes et organismes</json:string></inist></categories><publicationDate>2009</publicationDate><copyrightDate>2009</copyrightDate><doi><json:string>10.1111/j.1365-2966.2008.14283.x</json:string></doi><id>D9AEBE63A9901DB9EBE267E3263BC9DC7B158841</id><score>1</score><fulltext><json:item><extension>pdf</extension><original>true</original><mimetype>application/pdf</mimetype><uri>https://api.istex.fr/document/D9AEBE63A9901DB9EBE267E3263BC9DC7B158841/fulltext/pdf</uri></json:item><json:item><extension>zip</extension><original>false</original><mimetype>application/zip</mimetype><uri>https://api.istex.fr/document/D9AEBE63A9901DB9EBE267E3263BC9DC7B158841/fulltext/zip</uri></json:item><istex:fulltextTEI uri="https://api.istex.fr/document/D9AEBE63A9901DB9EBE267E3263BC9DC7B158841/fulltext/tei"><teiHeader><fileDesc><titleStmt><title level="a">Physical and chemical conditions in methanol maser selected hot cores and UCH ii regions</title></titleStmt><publicationStmt><authority>ISTEX</authority><publisher scheme="https://publisher-list.data.istex.fr">Blackwell Publishing Ltd</publisher><pubPlace>Oxford, UK</pubPlace><availability><licence><p>© 2009 The Authors. 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R.</forename><surname>Purcell</surname></persName><email>Cormac.Purcell@manchester.ac.uk</email><email>Cormac.Purcell@manchester.ac.uk</email><affiliation>University of Manchester, Jodrell Bank Observatory, Macclesfield, Cheshire SK11 9DL</affiliation><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation></affiliation></author><author xml:id="author-0001"><persName><forename type="first">S. N.</forename><surname>Longmore</surname></persName><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Harvard–Smithsonian Center For Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA</affiliation><affiliation>CSIRO Australia Telescope National Facility, PO Box 76, Epping, NSW 1710, Australia</affiliation></author><author xml:id="author-0002"><persName><forename type="first">M. G.</forename><surname>Burton</surname></persName><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation></author><author xml:id="author-0003"><persName><forename type="first">A. J.</forename><surname>Walsh</surname></persName><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Centre for Astronomy, James Cook University, Townsville, QLD 4811, Australia</affiliation></author><author xml:id="author-0004"><persName><forename type="first">V.</forename><surname>Minier</surname></persName><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Service d'Astrophysique, DAPNIA/DSM/CEA Saclay, 91191 Gif-sur-Yvette, France</affiliation><affiliation>AIM, Unité Mixte de Recherche, CEA-CNRS-Université Paris VII, UMR 7158, CEA/Saclay, 91191 Gif-sur-Yvette, France</affiliation></author><author xml:id="author-0005"><persName><forename type="first">M. R.</forename><surname>Cunningham</surname></persName><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation></author><author xml:id="author-0006"><persName><forename type="first">R.</forename><surname>Balasubramanyam</surname></persName><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Raman Research Institute, Sadashivanagar, Bangalore 560 080, India</affiliation></author><idno type="istex">D9AEBE63A9901DB9EBE267E3263BC9DC7B158841</idno><idno type="ark">ark:/67375/HXZ-7HBKJRZ7-G</idno><idno type="DOI">10.1111/j.1365-2966.2008.14283.x</idno></analytic><monogr><title level="j">Monthly Notices of the Royal Astronomical Society</title><title level="j" type="abbrev">Monthly Notices of the Royal Astronomical Society</title><idno type="pISSN">0035-8711</idno><idno type="eISSN">1365-2966</idno><idno type="publisher-id">mnras</idno><idno type="PublisherID-hwp">mnras</idno><imprint><publisher>Blackwell Publishing Ltd</publisher><pubPlace>Oxford, UK</pubPlace><date type="published" when="2009-03-21"></date><biblScope unit="volume">394</biblScope><biblScope unit="issue">1</biblScope><biblScope unit="page" from="323">323</biblScope><biblScope unit="page" to="339">339</biblScope></imprint></monogr></biblStruct></sourceDesc></fileDesc><profileDesc><creation><date>2009-03-07</date></creation><abstract><p>We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67 GHz methanol maser selected star-forming clumps observed by Purcell. In addition to the previously reported measurements of HCO+ (1–0), H13CO+ (1–0) and CH3CN (5–4) and (6–5), we used the Mopra antenna to detect emission lines of N2H+ (1–0), HCN (1–0) and HNC (1–0) towards 82/83 clumps (99 per cent), and CH3OH (2–1) towards 78/83 clumps (94 per cent). The molecular line data have been used to derive virial and local thermodynamic equilibrium masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between subsamples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6 GHz radio emission, indicating the presence of an Ultra-Compact H ii (UCH ii) region, and ‘isolated’ masers (without associated radio emission), and between clumps exhibiting CH3CN emission and those without. In particular, thermal CH3OH is found to be brighter and more abundant in UCH ii regions and in sources with detected CH3CN, and may constitute a crude molecular clock in single dish observations. Clumps associated with 8.6 GHz radio emission tend to be more massive and more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Hyper-Compact H ii or UCH ii region is the first visible tracer of star formation. The gas mass to submm/infrared luminosity relation for the combined sample was found to be L∝M0.68, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than, the mass of the stellar content. We also find that the mass of the hot core is correlated with the mass of the clump in which it is embedded.</p></abstract><textClass><keywords scheme="keyword"><list><head>keywords</head><item><term>surveys</term></item><item><term>stars: formation</term></item><item><term>stars: pre-main-sequence</term></item><item><term>ISM: abundances</term></item><item><term>ISM: molecules</term></item></list></keywords></textClass><textClass><keywords scheme="Journal Subject"><list><head></head><item><term>Papers</term></item></list></keywords></textClass></profileDesc><revisionDesc><change when="2009-03-07">Created</change><change when="2009-03-21">Published</change></revisionDesc></teiHeader></istex:fulltextTEI><json:item><extension>txt</extension><original>false</original><mimetype>text/plain</mimetype><uri>https://api.istex.fr/document/D9AEBE63A9901DB9EBE267E3263BC9DC7B158841/fulltext/txt</uri></json:item></fulltext><metadata><istex:metadataXml wicri:clean="corpus oup, element #text not found" wicri:toSee="no header"><istex:xmlDeclaration>version="1.0" encoding="utf-8"</istex:xmlDeclaration><istex:docType PUBLIC="-//NLM//DTD Journal Publishing DTD v2.3 20070202//EN" URI="journalpublishing.dtd" name="istex:docType"></istex:docType><istex:document><article article-type="research-article">
<front>
<journal-meta>
<journal-id journal-id-type="hwp">mnras</journal-id>
<journal-id journal-id-type="publisher-id">mnras</journal-id>
<journal-title>Monthly Notices of the Royal Astronomical Society</journal-title>
<abbrev-journal-title>Monthly Notices of the Royal Astronomical Society</abbrev-journal-title>
<issn pub-type="ppub">0035-8711</issn>
<issn pub-type="epub">1365-2966</issn>
<publisher>
<publisher-name>Blackwell Publishing Ltd</publisher-name>
<publisher-loc>Oxford, UK</publisher-loc>
</publisher>
</journal-meta>
<article-meta>
<article-id pub-id-type="doi">10.1111/j.1365-2966.2008.14283.x</article-id>
<article-categories>
<subj-group>
<subject>Papers</subject>
</subj-group>
</article-categories>
<title-group>
<article-title>Physical and chemical conditions in methanol maser selected hot cores and UCH <sc>ii</sc> regions</article-title>
</title-group>
<contrib-group>
<contrib contrib-type="author">
<name>
<surname>Purcell</surname>
<given-names>C. R.</given-names>
</name>
<xref ref-type="aff" rid="a1">1</xref><xref ref-type="aff" rid="a2">2</xref><xref ref-type="corresp" rid="c1">†</xref>
</contrib>
<contrib contrib-type="author">
<name>
<surname>Longmore</surname>
<given-names>S. N.</given-names>
</name>
<xref ref-type="aff" rid="a2">2</xref><xref ref-type="aff" rid="a3">3</xref><xref ref-type="aff" rid="a4">4</xref>
</contrib>
<contrib contrib-type="author">
<name>
<surname>Burton</surname>
<given-names>M. G.</given-names>
</name>
<xref ref-type="aff" rid="a2">2</xref>
</contrib>
<contrib contrib-type="author">
<name>
<surname>Walsh</surname>
<given-names>A. J.</given-names>
</name>
<xref ref-type="aff" rid="a2">2</xref><xref ref-type="aff" rid="a5">5</xref>
</contrib>
<contrib contrib-type="author">
<name>
<surname>Minier</surname>
<given-names>V.</given-names>
</name>
<xref ref-type="aff" rid="a2">2</xref><xref ref-type="aff" rid="a6">6</xref><xref ref-type="aff" rid="a7">7</xref>
</contrib>
<contrib contrib-type="author">
<name>
<surname>Cunningham</surname>
<given-names>M. R.</given-names>
</name>
<xref ref-type="aff" rid="a2">2</xref>
</contrib>
<contrib contrib-type="author">
<name>
<surname>Balasubramanyam</surname>
<given-names>R.</given-names>
</name>
<xref ref-type="aff" rid="a2">2</xref><xref ref-type="aff" rid="a8">8</xref>
</contrib>
</contrib-group>
<aff id="a1"><label>1</label>University of Manchester, Jodrell Bank Observatory, Macclesfield, Cheshire SK11 9DL</aff>
<aff id="a2"><label>2</label>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</aff>
<aff id="a3"><label>3</label>Harvard–Smithsonian Center For Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA</aff>
<aff id="a4"><label>4</label>CSIRO Australia Telescope National Facility, PO Box 76, Epping, NSW 1710, Australia</aff>
<aff id="a5"><label>5</label>Centre for Astronomy, James Cook University, Townsville, QLD 4811, Australia</aff>
<aff id="a6"><label>6</label>Service d'Astrophysique, DAPNIA/DSM/CEA Saclay, 91191 Gif-sur-Yvette, France</aff>
<aff id="a7"><label>7</label>AIM, Unité Mixte de Recherche, CEA-CNRS-Université Paris VII, UMR 7158, CEA/Saclay, 91191 Gif-sur-Yvette, France</aff>
<aff id="a8"><label>8</label>Raman Research Institute, Sadashivanagar, Bangalore 560 080, India</aff>
<author-notes>
<corresp id="c1">†E-mail: <email>Cormac.Purcell@manchester.ac.uk</email></corresp>
</author-notes>
<pub-date pub-type="ppub">
<day>21</day>
<month>03</month>
<year>2009</year>
</pub-date>
<pub-date pub-type="epub">
<day>07</day>
<month>03</month>
<year>2009</year>
</pub-date>
<volume>394</volume>
<issue>1</issue>
<fpage>323</fpage>
<lpage>339</lpage>
<history>
<date date-type="accepted"><day>20</day><month>11</month><year>2008</year></date>
<date date-type="received"><day>19</day><month>11</month><year>2008</year></date>
</history>
<copyright-statement>© 2009 The Authors. Journal compilation © 2009 RAS</copyright-statement>
<copyright-year>2009</copyright-year>
<abstract>
<p>We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67 GHz methanol maser selected star-forming clumps observed by Purcell. In addition to the previously reported measurements of HCO<sup>+</sup> (1–0), H<sup>13</sup>CO<sup>+</sup> (1–0) and CH<sub>3</sub>CN (5–4) and (6–5), we used the Mopra antenna to detect emission lines of N<sub>2</sub>H<sup>+</sup> (1–0), HCN (1–0) and HNC (1–0) towards 82/83 clumps (99 per cent), and CH<sub>3</sub>OH (2–1) towards 78/83 clumps (94 per cent).</p>
<p>The molecular line data have been used to derive virial and local thermodynamic equilibrium masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between subsamples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6 GHz radio emission, indicating the presence of an Ultra-Compact H <sc>ii</sc> (UCH <sc>ii</sc>) region, and ‘isolated’ masers (without associated radio emission), and between clumps exhibiting CH<sub>3</sub>CN emission and those without. In particular, thermal CH<sub>3</sub>OH is found to be brighter and more abundant in UCH <sc>ii</sc> regions and in sources with detected CH<sub>3</sub>CN, and may constitute a crude molecular clock in single dish observations.</p>
<p>Clumps associated with 8.6 GHz radio emission tend to be more massive <italic>and</italic> more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Hyper-Compact H <sc>ii</sc> or UCH <sc>ii</sc> region is the first visible tracer of star formation.</p>
<p>The gas mass to submm/infrared luminosity relation for the combined sample was found to be <italic>L</italic>∝<italic>M</italic><sup>0.68</sup>, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than, the mass of the stellar content.</p>
<p>We also find that the mass of the hot core is correlated with the mass of the clump in which it is embedded.</p>
</abstract>
<kwd-group>
<kwd>surveys</kwd>
<kwd>stars: formation</kwd>
<kwd>stars: pre-main-sequence</kwd>
<kwd>ISM: abundances</kwd>
<kwd>ISM: molecules</kwd>
</kwd-group>
</article-meta>
</front>
<body>
<sec id="ss1">
<title>1 INTRODUCTION</title>
<p>Molecular emission is a powerful tool when used to investigate the physical and chemical conditions in hot cores. Transitions requiring different temperatures and densities for excitation constitute an excellent probe of physical structure. Because the chemical properties of hot cores vary with time, the relative molecular abundances can also be used as indicators of evolution.</p>
<p>In the preceding decades, representative line surveys of a limited sample of cores have begun to identify the molecules most suited to investigating the process of massive star formation. To date, these have been restricted to a few objects: Orion-KL (<xref ref-type="bibr" rid="b3">Blake et al. 1986</xref>; <xref ref-type="bibr" rid="b47">Turner 1989</xref>; <xref ref-type="bibr" rid="b52">Ziurys & McGonagle 1993</xref>; <xref ref-type="bibr" rid="b41">Schilke et al. 1997</xref>), Sgr-B2 (<xref ref-type="bibr" rid="b13">Cummins, Linke & Thaddeus 1986</xref>; <xref ref-type="bibr" rid="b47">Turner 1989</xref>; <xref ref-type="bibr" rid="b44">Sutton et al. 1991</xref>), G34.3+0.15 (<xref ref-type="bibr" rid="b29">MacDonald et al. 1996</xref>; <xref ref-type="bibr" rid="b22">Kim et al. 2000</xref>, <xref ref-type="bibr" rid="b23">2001</xref>), IRAS 17470−2853 (also known as G5.89−0.39, <xref ref-type="bibr" rid="b24">Kim, Balasubramanyam & Burton 2002</xref>) and G305.2+0.2 (<xref ref-type="bibr" rid="b48">Walsh & Burton 2006</xref>).</p>
<p>Chemical models, including both gas-phase and grain-surface reactions, have been developed for specific objects whose physical structure and conditions are well known, e.g. G34.3+0.15 (<xref ref-type="bibr" rid="b32">Millar, MacDonald & Gibb 1997</xref>; <xref ref-type="bibr" rid="b45">Thompson, MacDonald & Millar 1999</xref>), as well as for general cases, e.g. <xref ref-type="bibr" rid="b39">Rodgers & Charnley (2003)</xref>. These and earlier models are limited by the computational power available at the time and consider either time-dependent chemistry at a test position or a ‘snapshot’ of spatial abundances at a single time. Advances in computational power have recently allowed the development of models which consider both time- <italic>and</italic> space-dependent chemistry simultaneously, e.g. the model of AFGL 2591 by <xref ref-type="bibr" rid="b15">Doty et al. (2002)</xref>. Today it is the lack of observational constraints which restrict our understanding of chemistry in young massive stellar objects. The priority in the coming years must be to assemble a large sample of molecular abundance measurements towards a broad range of massive star-forming regions, at different stages of evolution and with luminosities and masses from B3 to O5 type massive stars.</p>
<p>In this paper, we present the remaining results of the Mopra ‘Hot Molecular Cores’ (HMC) survey towards 6.67 GHz methanol (CH<sub>3</sub>OH) maser selected targets. The initial results were presented in <xref ref-type="bibr" rid="b37">Purcell</xref> (2006, hereafter Paper 1), wherein we reported the detection of the hot-core species methyl-cyanide (CH<sub>3</sub>CN) and the observations of the formyl ion (HCO<sup>+</sup>), and its isotopomer (H<sup>13</sup>CO<sup>+</sup>). Here, we expand the analysis to cover thermal transitions of methanol (CH<sub>3</sub>OH), carbonmonoxide (<sup>13</sup>CO), diazenylium (N<sub>2</sub>H<sup>+</sup>), hydrocyanic acid (HCN) and hydroisocyanic acid (HNC). We analyse the full complement of molecules for evidence of evolution.</p>
<p>Note: Tables S3–S9 are available in electronic form only along with additional supporting text and figures (see Appendix A).</p>
</sec>
<sec id="ss2">
<title>2 OBSERVATIONS</title>
<p>Observations were conducted on the Mopra millimetre wave telescope over five years 2000–2004 during the winter observing season, June to September. We observed the transitions CH<sub>3</sub>OH (2–1), <sup>13</sup>CO (1–0), N<sub>2</sub>H<sup>+</sup> (1–0), HCN (1–0) and HNC (1–0) as single pointings, targeted at the maser sites. Electronic properties of the molecules and transitions are presented in <xref ref-type="table" rid="t1">Table 1</xref>. Five sources do not contain methanol masers but were detected as discrete clumps at 1.2 mm, adjacent to clumps containing CH<sub>3</sub>OH masers (<xref ref-type="bibr" rid="b19">Hill et al. 2005</xref>). These ‘mm-only’ sources were observed as potential precursors to the hot-core phase of massive star formation. In each case, we targeted the peak of the 450 μm continuum emission later detected by <xref ref-type="bibr" rid="b50">Walsh et al. (2003)</xref>. The signal from the receiver was processed in an autocorrelator back-end configured to have a bandwidth of 64 MHz, split into 1024 channels. At 3-mm wavelengths, this delivers a velocity resolution of ∼0.2 km s<sup>−1</sup> over a usable velocity interval of ∼130 km s<sup>−1</sup>.</p>
<table-wrap id="t1">
<label>Table 1</label>
<caption><p>Details of observed transitions.</p></caption>
<table frame="hsides" rules="groups">
<thead>
<tr>
<td>Species</td>
<td>Transition</td>
<td>Frequency (GHz)</td>
<td><italic>E</italic> <sub>u</sub>/<italic>k</italic><sup>κ</sup> (K)</td>
<td><italic>A</italic> <sub>ul</sub> (s<sup>−1</sup>)</td>
<td><italic>g</italic> <sub>u</sub></td>
</tr>
</thead>
<tbody>
<tr>
<td rowspan="4" valign="top">CH<sub>3</sub>OH <sup>α</sup></td>
<td>2<sub>(−1,2)</sub>→ 1<sub>(−1,1)</sub> E</td>
<td>96.739 390</td>
<td>4.642</td>
<td>2.495 × 10<sup>−6</sup></td>
<td>5</td>
</tr>
<tr>
<td>2<sub>(0,2)</sub> → 1<sub>(0,1)</sub> A+</td>
<td>96.741 420</td>
<td>6.963</td>
<td>3.327 × 10<sup>−6</sup></td>
<td>5</td>
</tr>
<tr>
<td>2<sub>(0,2)</sub> → 1<sub>(0,1)</sub> E</td>
<td>96.744 580</td>
<td>12.188</td>
<td>3.327 × 10<sup>−6</sup></td>
<td>5</td>
</tr>
<tr>
<td>2<sub>(1,1)</sub> → 1<sub>(1,0)</sub> E</td>
<td>96.755 510</td>
<td>20.108</td>
<td>2.496 × 10<sup>−6</sup></td>
<td>5</td>
</tr>
<tr>
<td><sup>13</sup>CO <sup>β</sup></td>
<td>1 → 0</td>
<td>110.210 353</td>
<td>5.288</td>
<td>6.389 × 10<sup>−8</sup></td>
<td>3</td>
</tr>
<tr>
<td>HNC <sup>β</sup></td>
<td>1 → 0</td>
<td>90.663 593</td>
<td>4.350</td>
<td>2.709 × 10<sup>−5</sup></td>
<td>3</td>
</tr>
<tr>
<td rowspan="3" valign="top">HCN <sup>γ</sup></td>
<td>1<sub>1</sub>→ 0<sub>1</sub></td>
<td>88.630 416</td>
<td>4.253</td>
<td>2.444 × 10<sup>−5</sup></td>
<td>3</td>
</tr>
<tr>
<td>1<sub>2</sub>→ 0<sub>1</sub></td>
<td>88.631 847</td>
<td>–</td>
<td>–</td>
<td>5</td>
</tr>
<tr>
<td>1<sub>0</sub>→ 0<sub>1</sub></td>
<td>88.633 936</td>
<td>–</td>
<td>–</td>
<td>1</td>
</tr>
<tr>
<td rowspan="3" valign="top">N<sub>2</sub>H<sup>+</sup> <sup>γ</sup></td>
<td>1<sub>1</sub>→ 0<sub>1</sub></td>
<td>93.171 880</td>
<td>4.471</td>
<td>3.654 × 10<sup>−5</sup></td>
<td>9</td>
</tr>
<tr>
<td>1<sub>2</sub>→ 0<sub>1</sub></td>
<td>93.173 700</td>
<td>–</td>
<td>–</td>
<td>15</td>
</tr>
<tr>
<td>1<sub>0</sub>→ 0<sub>1</sub></td>
<td>93.176 130</td>
<td>–</td>
<td>–</td>
<td>3</td>
</tr>
</tbody>
</table>
<table-wrap-foot>
<fn id="t1_note71">
<p><sup>α</sup> <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu13.gif"></inline-graphic></inline-formula> quantum numbers.</p>
</fn>
<fn id="t1_note72">
<p><sup>β</sup> <italic>J</italic>→<italic>J</italic>′ quantum numbers.</p>
</fn>
<fn id="t1_note73">
<p><sup>γ</sup> <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu14.gif"></inline-graphic></inline-formula> quantum numbers.</p>
</fn>
<fn id="t1_note74">
<p><sup>κ</sup>Energies of the upper levels relative to the <italic>J<sub>K</sub></italic>= 1<sub>−1</sub> level for E-type CH<sub>3</sub>OH and relative to the <italic>J<sub>K</sub></italic>= 0<sub>0</sub> level for A-type CH<sub>3</sub>OH.</p>
</fn>
</table-wrap-foot>
</table-wrap>
<p>All observations were performed in a position switching mode by integrating for equal times on the science object and an emission-free reference position. Initial reduction of the raw data was performed using the <sc>spc</sc><xref ref-type="fn" rid="fn1">1</xref> package. Polynomial baselines were subtracted from the line-free channels using the <sc>xs</sc><xref ref-type="fn" rid="fn2">2</xref> software package. The <sc>class</sc><xref ref-type="fn" rid="fn3">3</xref> spectral line analysis software was used for all further processing and analysis, including baseline subtraction and fitting Gaussian line profiles. For further details on the observing setup and data analysis, please refer to <xref ref-type="sec" rid="ss3">section 3</xref> of Paper 1.</p>
</sec>
<sec id="ss3">
<title>3 RESULTS</title>
<p>Thermal molecular emission was detected in 82/83 targets, with the exception being the source G10.10−0.72. In the interim period, this source has been reobserved at 6.67 GHz using the 30-m Ceduna radio telescope (University of Tasmania) and no CH<sub>3</sub>OH maser was detected (Ellingsen, private communication). This source is thus most likely a noise artefact in the <xref ref-type="bibr" rid="b49">Walsh et al. (1998)</xref> survey and is discounted in the remainder of this work. Emission from <sup>13</sup>CO(1–0), N<sub>2</sub>H<sup>+</sup> (1–0), HNC (1–0) and HCN (1–0) was detected towards the remaining 82 sources, while CH<sub>3</sub>OH (2–1) emission was detected towards 79/82 sources (96 per cent). The 3σ noise limit on the spectra where no lines were detected is recorded in <xref ref-type="table" rid="t2">Table 2</xref>. <xref ref-type="fig" rid="f1">Fig. 1</xref> presents sample spectra for the source G323.74−0.26. The vertical dashed line marks the systemic velocity of the gas, measured from the optically thin H<sup>13</sup>CO<sup>+</sup> line (Paper 1). In the case of CH<sub>3</sub>OH, the rest frame is centred on the A+ transition; and for N<sub>2</sub>H<sup>+</sup> and HCN, the rest frame is centred on the <italic>J</italic><sub>F</sub>= 1<sub>2</sub>→ 0<sub>1</sub> E transition. Plots of all molecular spectra for all 83 observed sources (including G10.10−0.72) are available as additional online material (see Appendix A).</p>
<table-wrap id="t2">
<label>Table 2</label>
<caption><p>Noise detection limit on the spectra with no emission.</p></caption>
<table frame="hsides" rules="groups">
<thead>
<tr>
<td rowspan="2" valign="top">Source</td>
<td colspan="5" align="center"><italic>T</italic> <sub>MB</sub> 3σ detection limit (mK)</td>
</tr>
<tr>
<td>CH<sub>3</sub>OH</td>
<td><sup>13</sup>CO</td>
<td>N<sub>2</sub>H<sup>+</sup></td>
<td>HNC</td>
<td>HCN</td>
</tr>
</thead>
<tbody>
<tr>
<td>G0.32−0.20</td>
<td>230</td>
<td>–</td>
<td>–</td>
<td>–</td>
<td>–</td>
</tr>
<tr>
<td>G6.61−0.08</td>
<td>104</td>
<td>–</td>
<td>–</td>
<td>–</td>
<td>–</td>
</tr>
<tr>
<td>G10.10−0.72<sup>α</sup></td>
<td>69</td>
<td>414</td>
<td>142</td>
<td>127</td>
<td>163</td>
</tr>
<tr>
<td>G30.78+0.23</td>
<td>95</td>
<td>–</td>
<td>–</td>
<td>–</td>
<td>–</td>
</tr>
</tbody>
</table>
<table-wrap-foot>
<fn id="t2_note31">
<p><sup>α</sup>Walsh et al. (1997) originally detected a 6.67 GHz methanol towards G10.10−0.72. Subsequent observations (Ellingsen, private communication) have failed to detect the maser emission, hence we discard the source from our sample in the remainder of this work.</p>
</fn>
</table-wrap-foot>
</table-wrap>
<fig position="float" id="f1">
<label>Figure 1</label>
<caption>
<p>Sample spectra for the maser site associated with G323.74−0.26. The vertical dashed line marks the systemic velocity of the source, measured from the peak of the optically thin H<sup>13</sup>CO<sup>+</sup> line profile (see Purcell 2006). Vertical dotted lines mark the expected velocity of other transitions of the same molecule. Individual Gaussian fits to the line profiles are plotted by a dotted grey line and the cumulative fit by a solid grey line. We did not attempt to fit Gaussian to HCN lines as the majority of profiles have complex shapes.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f1.gif"></graphic>
</fig>
<sec id="ss3-1">
<title>3.1 Line profile parameters</title>
<p>As with CH<sub>3</sub>CN in Paper 1, we fit the multiple transitions of CH<sub>3</sub>OH (2–1) simultaneously with four Gaussian lines having equal full width at half-maximum (FWHM) linewidths and whose line-centre separations were fixed to the theoretical values. The weakest component in the spectrum, the <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu1.gif"></inline-graphic></inline-formula> E line, was detected towards 50 per cent of sources. Individual line profiles are generally well fit by single Gaussian, and the average linewidth is 4.7 km s<sup>−1</sup>. The CH<sub>3</sub>OH line profiles of five sources (G0.21+0.00, G0.26+0.01, G6.54−0.11, G10.30−0.15 and G15.03−0.71) were too confused or too weak to be successfully fit with Gaussian. The parameters of the Gaussian fits to the remaining 74 spectra are presented in Table S3.</p>
<p>The N<sub>2</sub>H<sup>+</sup> (1–0) transition is split into seven hyperfine lines as illustrated in <xref ref-type="fig" rid="f2">Fig. 2</xref> (adapted from <xref ref-type="bibr" rid="b6">Caselli, Myers & Thaddeus 1995</xref>). The relatively broad linewidths (>2 km s<sup>−1</sup>) observed towards massive star-forming regions result in the seven components blending into three groups with roughly Gaussian shapes. Consequently, we have characterized the profiles using two analysis schemes. In the first instance, we used the hyperfine structure fitting (HFS) routine in the <sc>class</sc> software to simultaneously fit the N<sub>2</sub>H<sup>+</sup> profiles with seven Gaussian. Free parameters were the velocity (<italic>V</italic><sub>LSR</sub>), common linewidth, excitation temperature and optical depth. In the case of severely blended lines (as presented in this work), the fits are not well constrained and values for the optical depth and excitation temperature contain substantial uncertainties. Using this method, we find the mean N<sub>2</sub>H<sup>+</sup> linewidth of the sample is 3.0 km s<sup>−1</sup>. Because of the uncertainty introduced by the blending, <sc>class</sc> assigned an erroneous value of τ= 0.1 to all sources, rendering the resultant excitation temperatures meaningless. As an alternative measure of optical depth and integrated intensity, we also fit the three blended groups individually with single Gaussian. In total, we successfully fit Gaussian to 79 spectra. The N<sub>2</sub>H<sup>+</sup> emission in the source G0.26+0.01, G6.54−0.11 and G6.61−0.8 was too weak or too confused to achieve a reliable fit. We discuss the optical depth and column densities derived from the N<sub>2</sub>H<sup>+</sup> spectra in <xref ref-type="sec" rid="ss4-4">Section 4.4</xref>. The results of the <sc>class</sc> HFS fits and the parameters of the Gaussian fits are presented in Table S4.</p>
<fig position="float" id="f2">
<label>Figure 2</label>
<caption>
<p>Spectrum of N<sub>2</sub>H<sup>+</sup> (1–0) towards a dark cloud showing the seven hyperfine components (<xref ref-type="bibr" rid="b6">Caselli et al. 1995</xref>). Linewidths towards massive star-forming regions are four to five times greater, resulting in the hyperfine components blending into three groups, designated 1–3 in this work (dashed lines).</p>
</caption>
<graphic xlink:href="mnras0394-0323-f2.gif"></graphic>
</fig>
<p>Line profiles of HNC are observed to be similar to HCO<sup>+</sup>. The critical densities of the two molecules are approximately the same at ∼3 × 10<sup>5</sup> cm<sup>−3</sup> (<xref ref-type="bibr" rid="b42">Schilke et al. 1992</xref>) and it is likely that they trace the same gas. Both molecules exhibit high-velocity line wings. Self-absorbed HNC line profiles are common in the sample and we referred to the low optical depth H<sup>13</sup>CO<sup>+</sup> and N<sub>2</sub>H<sup>+</sup> line profiles when attempting to distinguish between self-absorption and multiple clouds along the line of sight. We fit self-absorbed profiles with a single Gaussian by masking off the absorption dip (see Paper 1; <xref ref-type="fig" rid="f3">Fig. 3</xref>). We also measured the line intensity by integrating under the line between velocity limits bracketing the positive line emission and we calculated an equivalent linewidth from
<disp-formula id="m1">
<label>1</label>
<graphic xlink:href="mnras0394-0323-m1.gif"></graphic>
</disp-formula>
In <xref ref-type="disp-formula" rid="m1">equation (1)</xref>, <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu2.gif"></inline-graphic></inline-formula> is the integrated intensity and <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu3.gif"></inline-graphic></inline-formula> is the peak main-beam brightness temperature. An HNC emission from the sources G0.21+0.00 and G0.26+0.01 was not measured as the line profiles were too confused. In total, 79 spectra were successfully fit with Gaussian, while 80 were measured by integrating under the line profile. The results of the fits and integrated intensity measurements are presented in Table S5.</p>
<fig position="float" id="f3">
<label>Figure 3</label>
<caption>
<p>Distributions of FWHM linewidths from the Gaussian fits to the spectra. The HNC and <sup>13</sup>CO line profiles were fit with single Gaussian components. CH<sub>3</sub>OH was fit with multiple Gaussian, where the linewidths were constrained to have the same value. The N<sub>2</sub>H<sup>+</sup> spectrum was fit using the HFS method in <sc>class</sc>. Seven Gaussian components with a common linewidth were fit to the three groups of blended lines. The value plotted here is the linewidth of an individual component. The mean CH<sub>3</sub>OH, HNC and <sup>13</sup>CO linewidth are all at approximately 4.5 km s<sup>−1</sup>, while the mean N<sub>2</sub>H<sup>+</sup> linewidth is at 3.0 km s<sup>−1</sup>.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f3.gif"></graphic>
</fig>
<p>The <sup>13</sup>CO line profiles often exhibit multiple components spread over the 64 MHz bandpass. Usually, a single bright component is present at the systemic velocity of the source, which is easily fit with a single Gaussian. Occasionally, several smaller components are crowded around the main line and these are fit as blended Gaussian lines. In general, the <sup>13</sup>CO profiles appear moderately optically thick but exhibit less self-absorption than HNC or HCO<sup>+</sup>. We again measure the integrated intensity and peak brightness temperature directly from the raw spectrum. The parameters of the fits and integrated intensity measurement of all 82 spectra are presented in Table S6.</p>
<p>HCN exhibited highly confused line profiles, with few sources exhibiting the expected three-component hyperfine structure. We attribute the complex profiles to a combination of high optical depths, multiple clouds along the line of sight and low level emission from broad-line wings tracing outflows. It is possible that some spectra may be contaminated by emission from the reference position, as HCN appears to be ubiquitous near the Galactic plane.</p>
<p>Like HCO<sup>+</sup>, high-velocity outflow wings are common in the spectra of HNC and, to a lesser extent, <sup>13</sup>CO. We fit broad Gaussian to these wings simultaneously with the main component, and their parameters are presented in Tables S5 and S6 marked with a ‘<italic>w</italic>’. The Gaussian parameters for blended lines are marked with a ‘<italic>b</italic>’ in the same tables.</p>
</sec>
<sec id="ss3-2">
<title>3.2 Linewidths</title>
<p><xref ref-type="fig" rid="f3">Fig. 3</xref> shows the distributions of FWHM linewidths for the four molecules which were fit with Gaussian. The median CH<sub>3</sub>OH, HNC and <sup>13</sup>CO linewidths are all at approximately 4.5 km s<sup>−1</sup>, while the median N<sub>2</sub>H<sup>+</sup> linewidth is uniformly lower at 3.0 km s<sup>−1</sup>. Kolmogorov–Smirnov (KS) tests yield probabilities above 50 per cent that the CH<sub>3</sub>OH, HNC and <sup>13</sup>CO linewidth distributions are similar, implying that these lines are emitted from gas with similar dynamics. In contrast, the N<sub>2</sub>H<sup>+</sup> distribution is significantly different, yielding probabilities of less than 0.0001 per cent in each case.</p>
<p>For a temperature of 58 K (see <xref ref-type="sec" rid="ss4-1">Section 4.1</xref>), the median intrinsic thermal linewidths for the molecules are as follows: CH<sub>3</sub>OH: 0.29 km s<sup>−1</sup>, <sup>13</sup>CO: 0.31 km s<sup>−1</sup>, HNC: 0.32 km s<sup>−1</sup> and N<sub>2</sub>H<sup>+</sup>: 0.31 km s<sup>−1</sup>. These values represent only 7 to 10 per cent of the measured linewidths. A non-thermal line broadening may arise from bulk motions, rotation or turbulence, which are difficult to distinguish at the resolutions of these observations. We can conclude that the N<sub>2</sub>H<sup>+</sup> emission emanates from a more quiescent region than the other molecules.</p>
</sec>
</sec>
<sec id="ss4">
<title>4 DERIVED PARAMETERS</title>
<p>We present the general physical parameters of the sources and the excitation temperatures, column densities and optical depths derived from the line profiles. The frequencies and electronic constants used in the calculations are collected in <xref ref-type="table" rid="t1">Table 1</xref>.</p>
<sec id="ss4-1">
<title>4.1 Gas mass and virial mass</title>
<p>Gas masses have been calculated from the 1.2-mm continuum flux density, taken from the work of <xref ref-type="bibr" rid="b19">Hill et al. (2005)</xref>. For a 1.2-mm integrated continuum flux density, <italic>S</italic><sub>ν</sub> the mass of gas is given by <xref ref-type="bibr" rid="b18">Hildebrand (1983)</xref>:
<disp-formula id="m2">
<label>2</label>
<graphic xlink:href="mnras0394-0323-m2.gif"></graphic>
</disp-formula>
where <italic>D</italic> is the distance to the source in metres, κ<sub>d</sub> is the mass absorption coefficient per unit mass of dust, <italic>B</italic><sub>ν</sub> is the Planck function for a blackbody of temperature <italic>T</italic><sub>dust</sub> and <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu4.gif"></inline-graphic></inline-formula> is the dust to gas mass ratio. We have adopted a value of 0.1 m<sup>2</sup> kg<sup>−1</sup> for κ<sub>d</sub> at 1.2 mm and assumed a dust to gas mass ratio of 0.01 (<xref ref-type="bibr" rid="b34">Ossenkopf & Henning 1994</xref>). The distance has been estimated from the <italic>V</italic><sub>LSR</sub> of the source using the Galactic rotation curve of <xref ref-type="bibr" rid="b5">Brand & Blitz (1993)</xref> (see Paper 1). Values for <italic>T</italic><sub>dust</sub> were taken from the cold component of the grey body fit to the spectral energy distribution (SED) presented in Paper 1. For sources where we had insufficient data to perform a fit, we adopted the average value of <italic>T</italic><sub>dust</sub>= 58 K, determined from 65 good fits. We note that these temperatures are not well constrained by the sparsely sampled SED and are generally higher than the values of 20 to 30 K adopted in the literature. For dust temperatures of 20 and 30 K, the gas mass derived here will be underestimated on average by factors of 3.6 and 2.1, respectively. Values for <italic>M</italic><sub>gas</sub> were derived for 66 sources and are recorded in column six of Table S7.</p>
<p>We estimated virial masses from the CH<sub>3</sub>OH, N<sub>2</sub>H<sup>+</sup>, <sup>13</sup>CO and HNC line profiles following the approach by <xref ref-type="bibr" rid="b30">MacLaren, Richardson & Wolfendale (1988)</xref>. In the simplest case, neglecting support by magnetic fields or internal heating sources, the total mass of a simple spherical system is given by
<disp-formula id="m3">
<label>3</label>
<graphic xlink:href="mnras0394-0323-m3.gif"></graphic>
</disp-formula>
where σ is the full three-dimensional velocity dispersion, <italic>r</italic> is the dust-radius of the cloud in pc, Δ<italic>v</italic> is the FWHM linewidth of the molecular line in km s<sup>−1</sup> and <italic>G</italic> is the gravitational constant in m<sup>3</sup> kg<sup>−1</sup> s<sup>−2</sup>. The virial mass is the minimum mass required in order for a cloud to be gravitationally bound, i.e. the cloud is bound if <italic>M</italic><sub>gas</sub>/<italic>M</italic><sub>vir</sub> > 1. If magnetic field support is important, the virial mass may be overestimated by up to a factor of 2 (<xref ref-type="bibr" rid="b30">MacLaren et al. 1988</xref>). Observationally, the linewidths may be artificially broadened due to the blending of multiple components along the line of sight, or in optically thick lines, due to radiative transfer effects. Enhanced linewidths will cause us to overestimate the virial mass by an unknown amount and the values quoted here should be considered upper limits. Columns 7–10 of Table S7 present the virial masses derived from the four molecules.</p>
<p>The distributions of <italic>M</italic><sub>gas</sub> and <italic>M</italic><sub>vir</sub> are illustrated in <xref ref-type="fig" rid="f4">Fig. 4</xref>. Optically thin N<sub>2</sub>H<sup>+</sup> and H<sup>13</sup>CO<sup>+</sup> produce similar estimates of the mean virial mass at 573 and 905 M<sub>⊙</sub>, respectively. <sup>13</sup>CO and CH<sub>3</sub>OH have the broadest linewidths, most likely due to optical depth effects or low-velocity outflows, leading to mean mass estimates of 1293 and 1305 M<sub>⊙</sub>, respectively. The values derived for <italic>M</italic><sub>gas</sub> and <italic>M</italic><sub>vir</sub> are consistent with having the same magnitude in the majority of clumps. The large uncertainties associated with the calculations prevent the interpretation of individual objects as gravitationally bound.</p>
<fig position="float" id="f4">
<label>Figure 4</label>
<caption>
<p>Histograms showing the distributions of the virial masses derived from N<sub>2</sub>H<sup>+</sup>, H<sup>13</sup>CO<sup>+</sup>, <sup>13</sup>CO and CH<sub>3</sub>OH, and the gas masses derived from the 1.2-mm continuum emission from the work of <xref ref-type="bibr" rid="b19">Hill et al. (2005)</xref>. Virial masses should be considered upper limits. The plot on the bottom right-hand panel shows the distribution of the <italic>M</italic><sub>gas</sub>/<italic>M</italic><sub>vir</sub> ratio derived from N<sub>2</sub>H<sup>+</sup>. The values of <italic>M</italic><sub>gas</sub> and <italic>M</italic><sub>vir</sub> are consistent with having the same order of magnitude, however, <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu11.gif"></inline-graphic></inline-formula> as <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu12.gif"></inline-graphic></inline-formula> is < Δ<italic>v</italic><sub>others</sub>.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f4.gif"></graphic>
</fig>
</sec>
<sec id="ss4-2">
<title>4.2 H<sub>2</sub> column and volume density</title>
<p>We estimated the H<sub>2</sub> column density, <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu5.gif"></inline-graphic></inline-formula>, from the gas mass using the following relation:
<disp-formula id="m4">
<label>4</label>
<graphic xlink:href="mnras0394-0323-m4.gif"></graphic>
</disp-formula>
where <italic>M</italic><sub>p</sub>= 1.673 × 10<sup>−27</sup> kg is the mass of a proton, <italic>f</italic><sub>He</sub>= 1.36 is a correction factor to account for Helium present in the interstellar medium (<xref ref-type="bibr" rid="b1">Allen 1973</xref>) and <italic>A</italic> is the projected surface area of the mm-continuum emission. Assuming a cylindrical source with a uniform density, the projected area is given by
<disp-formula id="m5">
<label>5</label>
<graphic xlink:href="mnras0394-0323-m5.gif"></graphic>
</disp-formula>
where <italic>r</italic> is the angular radius of the source in radians and <italic>D</italic> is the distance to the source in metres (1 pc = 3.085 68 × 10<sup>16</sup> m). In a similar manner, the volume density <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu6.gif"></inline-graphic></inline-formula> can be estimated by replacing <italic>A</italic> in <xref ref-type="disp-formula" rid="m4">equation (4)</xref> by <italic>V</italic> the projected volume of emitting gas
<disp-formula id="m6">
<label>6</label>
<graphic xlink:href="mnras0394-0323-m6.gif"></graphic>
</disp-formula>
<xref ref-type="disp-formula" rid="m6">Equation (6)</xref> assumes emission from a uniformly dense sphere of angular radius <italic>r</italic>. The radius of the dust continuum emission was measured directly from the 1.2-mm continuum images (provided courtesy of <xref ref-type="bibr" rid="b19">Hill et al. 2005</xref>) by fitting a Gaussian to the azimuthally averaged source profiles. We find values for the FWHM vary between 99 and 24 arcsec (the limiting resolution of the <sc>simba</sc> bolometer), and have a mean value of 34 arcsec. Derived values of <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu7.gif"></inline-graphic></inline-formula> and <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu8.gif"></inline-graphic></inline-formula> are presented in columns four and five of Table S7 and have means of 6.1 × 10<sup>22</sup> and 4.8 × 10<sup>4</sup> cm<sup>−3</sup>, respectively.</p>
<p>We urge caution in interpreting any value of the H<sub>2</sub> column density derived in this manner as the assumptions introduce large unknown errors. First, the mass absorption coefficient and dust to gas ratio used to calculate <italic>M</italic><sub>gas</sub> in <xref ref-type="disp-formula" rid="m2">equation (2)</xref> may vary considerably from source to source, leading to a corresponding error in the H<sub>2</sub> column density. The spatial distribution of H<sub>2</sub> is further unknown and additional errors are introduced by assuming a uniform column density over a spherical projected volume. For the six sources at the limit of <sc>simba</sc> resolution (24 arcsec) derived values should be considered beam-averaged lower limits. Also see <xref ref-type="sec" rid="ss6-3">Section 6.3</xref> for a discussion on beam dilution.</p>
</sec>
<sec id="ss4-3">
<title>4.3 CH<sub>3</sub>OH rotational temperature and column density</title>
<p>CH<sub>3</sub>OH is a slightly asymmetric rotor, resembling a symmetric top except for the O–H group, which is angled with respect to the principal axis. Internal hindered motion coupled with the rotation of the molecule results in a complicated rotational spectrum, organized into non-degenerate A and two-fold degenerate E levels, corresponding to different torsional symmetry. The A and E states have nuclear statistical weights of 2 and 1, respectively, thus their overall degeneracy is the same (<xref ref-type="bibr" rid="b4">Bockelee-Morvan et al. 1994</xref>).</p>
<p>A rotational temperature <italic>T</italic><sub>rot</sub> and column density <italic>N</italic> may be estimated from the ratios of the thermal CH<sub>3</sub>OH lines based on the assumptions of low optical depths and that the background temperature (<italic>T</italic><sub>b</sub>≈ 2.7 K) is small compared to the brightness temperature of the line. Under conditions of local thermodynamic equilibrium (LTE), <italic>all</italic> the observed energy levels can be described by a single excitation temperature <italic>T</italic><sub>ex</sub>, defined by the Boltzmann relation. Using the relation for rotational temperature (<xref ref-type="bibr" rid="b17">Goldsmith & Langer 1999</xref>),
<disp-formula id="m7">
<label>7</label>
<graphic xlink:href="mnras0394-0323-m7.gif"></graphic>
</disp-formula>
a rotation diagram may be drawn by plotting ln (<italic>N</italic><sub>u</sub>/<italic>g</italic><sub>u</sub>) versus <italic>E</italic><sub>u</sub>/<italic>k</italic>. The rotation temperature and column density in the upper level <italic>N</italic><sub>u</sub> are found from the slope and <italic>y</italic>-axis intercept, respectively, of a straight line fit to the data. The total column density <italic>N</italic> may then be found assuming LTE and a partition function <italic>Q</italic>(<italic>T</italic><sub>rot</sub>), which for CH<sub>3</sub>OH is given by (<xref ref-type="bibr" rid="b46">Townes & Schawlow 1955</xref>)
<disp-formula id="m8">
<label>8</label>
<graphic xlink:href="mnras0394-0323-m8.gif"></graphic>
</disp-formula>
In <xref ref-type="disp-formula" rid="m8">equation (8)</xref>, we have approximated the partition function by a fit to the discrete values in the <xref ref-type="bibr" rid="b35">Pickett et al. (1998)</xref> spectral line catalogue. Note that values for <italic>E</italic><sub>u</sub> are referenced to different ground-state energy levels in A- and E-type CH<sub>3</sub>OH. The ground state for A-type is the <italic>J</italic><sub>K</sub>= 0<sub>0</sub> level, while for E-type CH<sub>3</sub>OH is the <italic>J<sub>K</sub></italic>= 1<sub>−1</sub> level (<xref ref-type="bibr" rid="b31">Menten et al. 1988</xref>). This is commonly overlooked in the literature as all values quoted in <xref ref-type="bibr" rid="b35">Pickett et al. (1998)</xref> are referenced to the <italic>J<sub>K</sub></italic>= 0<sub>0</sub> level, resulting in incorrect rotation diagrams. The correct values of <italic>E</italic><sub>u</sub>/<italic>k</italic> are cited in <xref ref-type="table" rid="t1">Table 1</xref>.</p>
<p>An example of a CH<sub>3</sub>OH rotation diagram is shown in <xref ref-type="fig" rid="f5">Fig. 5</xref> (left-hand panel). The data point at <italic>E</italic><sub>u</sub>/<italic>k</italic>≈ 7 K corresponds to the A+ line, while the other three points correspond to E transitions. In general, where we detect three E-type lines, they are well fit with a straight line, implying that our assumption of optically thin conditions is valid.</p>
<fig position="float" id="f5">
<label>Figure 5</label>
<caption>
<p>Left-hand panel: rotational diagram of CH<sub>3</sub>OH for the source G10.32−0.16, made using the 96.7 GHz lines. The point at <italic>E</italic><sub>u</sub>≈ 7 K corresponds to the 2<sub>0</sub>→ 1<sub>0</sub> A+ transition. The three other points are due to transitions of E-type methanol. The solid line and values for <italic>T</italic><sub>rot</sub> and <italic>N</italic> correspond to a fit through the <italic>E</italic> transitions only. The dashed line represents a fit to all transitions. Right-hand panel: distributions of rotational temperatures derived from CH<sub>3</sub>OH using both A- and E-type transitions (top panel) and E-type transitions only (bottom panel).</p>
</caption>
<graphic xlink:href="mnras0394-0323-f5.gif"></graphic>
</fig>
<p>A and E type methanol are oblate and prolate forms, respectively, of the asymmetric top and only interconvert on time-scales of ∼10 000 yr. It is not clear that the A/E ratio can be assumed equal to 1 and a rotation diagram including both A and E type transitions may not be valid. In <xref ref-type="fig" rid="f5">Fig. 5</xref> (left-hand panel), the solid line indicates a fit to the E-type transitions only, while the dotted line is a fit to both the A and E type. Rotation temperatures derived from only the E-transitions have a mean value of 6.67 K, similar to the mean temperature derived from all four transitions, which is 6.8 K. Rotational temperatures and column densities derived using both fitting methods are presented in Table S8 and the distributions of rotational temperatures are illustrated in <xref ref-type="fig" rid="f5">Fig. 5</xref> (right-hand panel).</p>
<p>Column densities are averaged over the telescope beam and likely underestimated. If the bulk of the CH<sub>3</sub>OH emission stems from within the bounds of a hot core then the typical size of the emitting region will be ∼0.1 pc (≈20 000 au, <xref ref-type="bibr" rid="b26">Kurtz et al. 2000</xref>). Columns three and four of Table S8 present the angular size and beam dilution factor calculated using this assumption. Columns seven and ten present the column densities corrected for this assumed dilution factor.</p>
<p>On inspection of the bulk of the rotation diagrams (available as an online supplement), we see a pattern emerge. The A-type transition has a consistently higher <italic>N</italic><sub>u</sub>/<italic>g</italic><sub>u</sub> compared to a straight line fit through the E-type transitions. This indicates a greater abundance of A-type CH<sub>3</sub>OH and an A/E ratio >1. Considering only sources in which all three E-type lines were detected (36 sources), and assuming equal excitation temperatures in A and E lines, we derive A/E abundance ratios ranging from 0.47 to 1.80, with a mean of 1.30. These values are comparable to values previously derived in the literature, e.g. <xref ref-type="bibr" rid="b31">Menten et al. (1988)</xref> who calculated A/E ratios between 1.3 and 2.0, using the same transitions, towards warm clumps in Orion-KL.</p>
</sec>
<sec id="ss4-4">
<title>4.4 N<sub>2</sub>H<sup>+</sup> optical depth and column density</title>
<p>Optical depths were derived from the blended hyperfine components of N<sub>2</sub>H<sup>+</sup>, which assume set ratios under optically thin conditions (see <xref ref-type="sec" rid="ss3-1">Section 3.1</xref>). Assuming the linewidths of the individual hyperfine components are all equal, the integrated intensities of the three blended groups should be in the ratio of 1 : 5 : 2 under optically thin conditions. The optical depth is calculated from the ratio of integrated intensities <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu9.gif"></inline-graphic></inline-formula> of any two groups using the following equation:
<disp-formula id="m9">
<label>9</label>
<graphic xlink:href="mnras0394-0323-m9.gif"></graphic>
</disp-formula>
where ‘<italic>a</italic>’ is the expected ratio of τ<sub>2</sub>/τ<sub>1</sub> under optically thin conditions. <xref ref-type="bibr" rid="b6">Caselli et al. (1995)</xref> report anomalous excitation of the <italic>F</italic><sub>1</sub>, <italic>F</italic>= 1, 0 → 1, 1 and 1, 2 → 1, 2 components (in our group 3), so we determine the optical depth solely from the intensity ratio of group 1/group 2. We find that 97 per cent of sources have optical depths below 1. <xref ref-type="fig" rid="f6">Fig. 6</xref> plots the distribution of the derived optical depths, which range from 0.035 to 2.43, with a mean of 0.35 and a standard deviation of 0.42. Values for the optical depth towards individual sources are presented in column two of Table S9.</p>
<fig position="float" id="f6">
<label>Figure 6</label>
<caption>
<p>Distribution of optical-depths derived from the ratios of the hyperfine groups of N<sub>2</sub>H<sup>+</sup>. The mean is 0.35 ± 0.42.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f6.gif"></graphic>
</fig>
<p>Column densities of N<sub>2</sub>H<sup>+</sup> can be found from the following standard formula (<xref ref-type="bibr" rid="b40">Rohlfs & Wilson 2004</xref>):
<disp-formula id="m10">
<label>10</label>
<graphic xlink:href="mnras0394-0323-m10.gif"></graphic>
</disp-formula>
where <italic>Q</italic>(<italic>T</italic><sub>ex</sub>) is the partition function, extrapolated from a fit to the values given by the (<xref ref-type="bibr" rid="b35">Pickett et al. 1998</xref>) spectral line catalogue:
<disp-formula id="m11">
<label>11</label>
<graphic xlink:href="mnras0394-0323-m11.gif"></graphic>
</disp-formula>
<xref ref-type="disp-formula" rid="m10">Equation (10)</xref> requires an estimate of the excitation temperature <italic>T</italic><sub>ex</sub>. Unfortunately, under optically thin conditions, the excitation temperature cannot be determined from the N<sub>2</sub>H<sup>+</sup> line profiles, so we assume a constant value of 10 K. For excitation temperatures of 5 and 20 K, the total column density should be multiplied by 0.8 and 1.6, respectively. Beam-averaged column densities derived using <xref ref-type="disp-formula" rid="m10">equation (10)</xref> are presented in column five of Table S9. In the literature, high-resolution maps show that N<sub>2</sub>H<sup>+</sup> traces the spatial extent of the dust closely (e.g. <xref ref-type="bibr" rid="b7">Caselli et al. 2002a</xref>; <xref ref-type="bibr" rid="b36">Pirogov et al. 2003</xref>). We have estimated the average FWHM of the dust emission for all sources observed by <xref ref-type="bibr" rid="b19">Hill et al. (2005)</xref> (see <xref ref-type="sec" rid="ss4-2">Section 4.2</xref>) and use these values to correct for beam dilution. Total column densities corrected by a factor of (θ<sub>dust</sub>/θ<sub>beam</sub>)<sup>2</sup> are presented in column six of Table S9.</p>
</sec>
<sec id="ss4-5">
<title>4.5 <sup>13</sup>CO, HNC and HCN column density</title>
<p>The HNC, HCN and <sup>13</sup>CO profiles frequently display asymmetries which may indicate that they are optically thick in general. Lacking an estimate for both the excitation temperature and the optical depth, we are unable to calculate reliable column densities.</p>
</sec>
<sec id="ss4-6">
<title>4.6 Abundances</title>
<p>The relative abundance <italic>X</italic> between two species may be found directly from the ratio of their volume densities. Assuming both molecules occupy the same volume of space, this is simply equivalent to the ratio of their column densities:
<disp-formula id="m12">
<label>12</label>
<graphic xlink:href="mnras0394-0323-m12.gif"></graphic>
</disp-formula>
If, however, one or both species subtend angles smaller than the beam, we must correct for their <italic>relative</italic> volume filling factor and <xref ref-type="disp-formula" rid="m12">equation (12)</xref> becomes
<disp-formula id="m13">
<label>13</label>
<graphic xlink:href="mnras0394-0323-m13.gif"></graphic>
</disp-formula>
assuming spherically symmetric emission with characteristic sizes θ<sub>1</sub> and θ<sub>2</sub>. For the single-position observations presented here, the volume filling factors are complete unknowns and constitute the greatest source of errors in the abundance. Additional errors stem from the calibration of the brightness scale (∼30 per cent) and the pointing accuracy of the telescope (∼20 per cent). Abundances relative to H<sub>2</sub> depend also on the assumptions made in calculating the gas mass from the 1.2-mm continuum emission and cannot be determined reliably.</p>
</sec>
</sec>
<sec id="ss5">
<title>5 LUMINOSITIES AND MASSES</title>
<p>The relationship between mass and luminosity provides clues to the star formation efficiency, stellar content and evolutionary state of the clumps. <xref ref-type="fig" rid="f7">Fig. 7</xref> (left-hand side) is an expanded version of fig. 2 in <xref ref-type="bibr" rid="b9">Cesaroni (2005)</xref>, showing the relationship between luminosity and gas mass for the 55 clumps whose properties could be determined. The luminosity has been measured from a two-component grey body fit to the SED, and is directly related to the stellar content of the clump. Submillimetre emission from assumed optically thin dust (<xref ref-type="bibr" rid="b19">Hill et al. 2005</xref>) has been used to derive the total gas mass of the clumps. In the plot, the thick line marks the expected luminosity from a minimum number of stars of mass <italic>M</italic><sub>*</sub>≤ 100 M<sub>⊙</sub>, assuming <italic>M</italic><sub>gas</sub>≡<italic>M</italic><sub>stars</sub>. <xref ref-type="bibr" rid="b9">Cesaroni (2005)</xref> divides a sample of 21 known hot cores<xref ref-type="fn" rid="fn4">4</xref> into ‘light’ and ‘heavy’ categories, based on their position above or below this division. ‘Light’ cores have gas masses below ∼100 M<sub>⊙</sub><italic>and</italic> fall above the thick line. Their high luminosity-to-mass ratio implies that the mass of their stellar content is greater than that of the gas. ‘Heavy’ cores have gas masses above ∼100 M<sub>⊙</sub> and fall below the thick line. The mass of their stellar content is comparable to, or less than, the mass of the gas. <xref ref-type="bibr" rid="b9">Cesaroni (2005)</xref> has interpreted the <italic>L</italic><sub>stars</sub>/<italic>M</italic><sub>gas</sub> ratio in the ‘light’ cores as suggesting only a single embedded massive star exists in these clumps. Conversely, the ‘heavy’ clumps are likely to contain multiple massive stars which have formed in clustered mode. We do not find any ‘light’ cores in our sample.</p>
<fig position="float" id="f7">
<label>Figure 7</label>
<caption>
<p>Left-hand panel: plot of luminosity versus dust + gas mass for 55 sources whose properties could be determined (after <xref ref-type="fig" rid="f2">Fig. 2</xref>, <xref ref-type="bibr" rid="b9">Cesaroni 2005</xref>). UCH <sc>ii</sc> regions are represented by solid triangles and isolated masers by crosses. The thick solid line traces the luminosity expected if <italic>M</italic><sub>gas</sub>≡<italic>M</italic><sub>stars</sub> in the clumps, and the stellar population is concentrated in a minimum number of stars with masses ≤ 100 M<sub>⊙</sub>. The dashed line is the luminosity expected from the most massive star in a <italic>M</italic><sub>gas</sub> cluster with stellar masses distributed according a Salpeter IMF (α= 2.35), and a fundamental upper stellar mass limit of 150 M<sub>⊙</sub> (<xref ref-type="bibr" rid="b51">Weidner & Kroupa 2004</xref>). We have assumed a mass–luminosity relation of <italic>L</italic>∝<italic>M</italic><sup>3</sup>. The dash–dotted line traces the expected luminosity from the same IMF assuming <italic>L</italic>∝<italic>M</italic><sup>4</sup> and the shaded region indicates intermediate values. The observed mass–luminosity relation is fit with a single power law of index 0.68 ± 0.09 (thin solid line). Right-hand panel: histograms of the distribution in the <italic>M</italic><sub>gas</sub>/<italic>L</italic><sub>stars</sub> values for UCH <sc>ii</sc> regions and isolated masers. A KS test cannot distinguish between them.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f7.gif"></graphic>
</fig>
<p>An alternative explanation for high <italic>L</italic><sub>stars</sub>/<italic>M</italic><sub>gas</sub> ratios in the ‘light’ sources might derive from the age of the star-forming regions. A massive young stellar object at a late stage of evolution could be expected to have dispersed a large fraction of its molecular gas, through the action of stellar winds, intense ultraviolet-radiation (UV-radiation) and outflows from young OB stars. This would result in a decrease in the measured gas mass without a corresponding decrease in the luminosity. In <xref ref-type="fig" rid="f7">Fig. 7</xref>, the points would move horizontally into the ‘light’ portion of the plot. We note that Cesaroni's light cores include Orion-KL and IRAS 20126+4014, neither of which are believed to be at a late stage of evolution. Orion is the closest massive star-forming region and is not clear if it is a representative of massive star-forming regions in general. The mass of IRAS 20126+4014 (7 ± 3 M<sub>⊙</sub>) has been measured assuming Keplerian rotation of a molecular disc (<xref ref-type="bibr" rid="b10">Cesaroni et al. 2005</xref>) and relates to the inner object only. On larger scales, the gas mass has been measured to be several hundred solar masses (e.g. <xref ref-type="bibr" rid="b16">Estalella et al. 1993</xref>), putting IRAS 20126+4014 in the ‘heavy’ category. We find that all of our maser selected sample have gas masses above ∼30 M<sub>⊙</sub> and <italic>all</italic> fall into the ‘heavy’ category. If the division between ‘light’ and ‘heavy’ cores is an indication of age, then all of our sample are at a relatively early stage of evolution, i.e. the powering massive young stellar objects have not yet dispersed their natal molecular clouds.</p>
<p>In <xref ref-type="fig" rid="f7">Fig. 7</xref>, the shaded region between the dashed lines represents the range of luminosities expected from the <italic>most massive</italic> star in a cluster where <italic>M</italic><sub>stars</sub>=<italic>M</italic><sub>gas</sub>, if the initial mass function (IMF) follows a Salpeter power law (α= 2.35) and assuming a fundamental upper stellar mass limit of <italic>M</italic><sub>*</sub>= 150 M<sub>⊙</sub> (<xref ref-type="bibr" rid="b51">Weidner & Kroupa 2004</xref>). The luminosity of the cluster will be dominated by the most massive star so this is a good approximation to the expected total luminosity. The range in expected luminosities derives from the uncertainty in the mass–luminosity relation for massive young stellar objects. This may be represented by a power law of the form <italic>L</italic><sub>*</sub>∝<italic>M</italic><sup>γ</sup><sub>*</sub>, where γ ranges between 3 and 4 in the literature for OB stars (e.g. <xref ref-type="bibr" rid="b27">Lada 1999</xref>; <xref ref-type="bibr" rid="b12">Churchwell 2002</xref>). Our maser selected sample exhibits a rough correlation between mass and luminosity, which may be fit with a single power law with index γ= 0.68 ± 0.08. This is a much shallower rise than observed in evolved stars, however, the luminosities of the most massive clusters are consistent with the 150 M<sub>⊙</sub> stellar cut-off (<xref ref-type="bibr" rid="b51">Weidner & Kroupa 2004</xref>), even in the more conservative limit of <italic>L</italic><sub>*</sub>∝<italic>M</italic><sup>3</sup><sub>*</sub>.</p>
<p>The derived gas mass is highly dependent on the temperature in <xref ref-type="disp-formula" rid="m2">equation (2)</xref>. The shallow power law might be explained if the average dust temperature increases with luminosity and this change is not reflected in the sparsely sampled grey body fits to the SED, from which we take our temperature estimates. If we assume the mass estimates are correct to within a factor of a few, then the main result is that the lower mass clumps are overluminous and the higher mass clumps are underluminous.</p>
<sec id="ss5-1">
<title>5.1 UCH <sc>ii</sc> regions: luminosity or age?</title>
<p>UCH <sc>ii</sc> regions (triangles) are, in general, more luminous <italic>and</italic> more massive than the isolated masers (crosses) in the sample. Individual power-law fits to the isolated masers and UCH <sc>ii</sc> subsamples yield mass–luminosity relations of <italic>L</italic><sub>bol | masers</sub>∝<italic>M</italic><sup>0.59±0.13</sup> and <italic>L</italic><sub>bol | UCH II</sub>∝<italic>M</italic><sup>0.64±0.18</sup>, respectively. These fits are consistent with the fit to the combined sample, within the errors, and a single power law can describe both distributions. <xref ref-type="fig" rid="f7">Fig. 7</xref> (right-hand panel) is a histogram of the <italic>M</italic><sub>gas</sub>/<italic>L</italic><sub>stars</sub> ratio for UCH <sc>ii</sc> regions (top panel) and isolated masers (bottom panel). As expected from the mass–luminosity plot, the distributions are similar and a KS test cannot distinguish between them, yielding a 76 per cent probability of the two subsamples being drawn from the same population. This result implies that for the same masses, the two samples have the same luminosities. Early in the evolution of a massive young stellar object, we would expect its luminosity to increase as it accretes matter and more mass is transferred from the envelope to the star. At the same time, the overall mass of gas and dust in the clump is expected to remain relatively constant, or indeed decrease, as winds from the young OB cluster break up the cloud. By this rationale, if our sample of UCH <sc>ii</sc> regions is older and more evolved, we would expect them to have a lower <italic>M</italic><sub>gas</sub>/<italic>L</italic><sub>stars</sub> ratio than the isolated masers. There are several possible reasons why we do not see a difference between the mass–luminosity distributions.</p>
<p>It is possible that systematic errors (which we discuss below) associated with our assumptions cause us to overestimate the mass of the clouds hosting the UCH <sc>ii</sc> regions relative to the clouds hosting the isolated masers. For there to be no difference between the mean masses of the subsamples, the isolated masers would need to have a dust temperature on average 36 K lower than the UCH <sc>ii</sc> regions. <xref ref-type="fig" rid="f8">Fig. 8</xref> (left-hand panel) illustrates the distributions of <italic>T</italic><sub>cold</sub> from the SED fits, for UCH <sc>ii</sc> regions (top panel) and isolated masers (bottom panel). While the UCH <sc>ii</sc> regions clearly have temperatures which are skewed to the high end, the difference in the means is only 4.3 K. The rotational temperatures calculated from CH<sub>3</sub>CN are likewise similar, however, we cannot rule out different average dust temperatures on the basis that CH<sub>3</sub>CN emission likely probes only the embedded ‘hot core’ and does not reflect the conditions in the extended clump. Thermal CH<sub>3</sub>OH emission is thought to derive from a more extended envelope and rotational temperatures calculated from this molecule appear to support very different temperatures in the two groups. <xref ref-type="fig" rid="f8">Fig. 8</xref> (right-hand panel) shows histograms of the CH<sub>3</sub>OH rotational temperature for the UCH <sc>ii</sc> regions (top panel) and isolated masers (bottom panel). A KS test returns a probability of 0.002 per cent of the distributions deriving from the same parent. We note that the mean values of 9.9 K for the UCH <sc>ii</sc> regions and 6.1 K for the isolated masers do not reflect the kinetic temperature as the CH<sub>3</sub>OH is likely subthermally excited (see <xref ref-type="bibr" rid="b17">Goldsmith & Langer 1999</xref>).</p>
<fig position="float" id="f8">
<label>Figure 8</label>
<caption>
<p>Left-hand panel: histograms showing the distributions of dust temperature taken from the cold component of the grey body fit to the SED. The top panel shows UCH <sc>ii</sc> regions, which have a mean <italic>T</italic><sub>cold</sub> of 62.4 K, while the bottom panel shows isolated masers, which have a mean <italic>T</italic><sub>cold</sub> of 58.1 K. Right-hand panel: histograms showing the distributions of CH<sub>3</sub>OH rotational temperature for UCH <sc>ii</sc> regions (top panel) and isolated masers (bottom panel). The mean values are 9.9 K for UCH <sc>ii</sc> regions and 6.1 K for isolated masers. A KS test establishes that the CH<sub>3</sub>OH rotational temperature distributions are significantly different, returning a probability of 0.002 per cent.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f8.gif"></graphic>
</fig>
<p>As an alternate estimate of the mass, we derived an ‘LTE–mass’ from the column density of N<sub>2</sub>H<sup>+</sup>. In our calculations, we assume (i) the N<sub>2</sub>H<sup>+</sup> emission closely traces the 1.2-mm dust continuum and has a flat density distribution within the FWHM of the clump and (ii) a fixed [H<sub>2</sub>] / [N<sub>2</sub>H<sup>+</sup>] abundance ratio of 1.7 × 10<sup>−10</sup> (<xref ref-type="bibr" rid="b36">Pirogov et al. 2003</xref>). The mass of gas may then be estimated from the column density using <xref ref-type="disp-formula" rid="m4">equation (4)</xref>. <xref ref-type="fig" rid="f9">Fig. 9</xref> (left-hand panel) is a mass–luminosity plot similar to <xref ref-type="fig" rid="f7">Fig. 7</xref>, except that the mass plotted on the <italic>x</italic>-axis is the LTE–mass calculated from N<sub>2</sub>H<sup>+</sup>. Although the scatter is greater, we clearly see the masses of the UCH <sc>ii</sc> regions and isolated masers are still weakly offset, with the UCH <sc>ii</sc> regions having greater masses on average, as in <xref ref-type="fig" rid="f7">Fig. 7</xref>. The distributions of the <italic>M</italic><sub>gas</sub>/<italic>L</italic><sub>stars</sub> ratios are illustrated in <xref ref-type="fig" rid="f9">Fig. 9</xref> (right-hand panel). A KS test returns a probability of 50 per cent that both distributions are derived from the same parent, confirming the original result.</p>
<fig position="float" id="f9">
<label>Figure 9</label>
<caption>
<p>As <xref ref-type="fig" rid="f7">Fig. 7</xref>, except in this case the mass is the LTE-mass derived from the column density of N<sub>2</sub>H<sup>+</sup> assuming a fixed abundance ratio of 1.7 × 10<sup>−10</sup> and that the N<sub>2</sub>H<sup>+</sup> emission closely follows the spacial distribution of the 1.2-mm continuum emission.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f9.gif"></graphic>
</fig>
<p>Another explanation for the offset in the masses derives from the sensitivity limit on the radio survey. In the preceding paper, we assumed the UCH <sc>ii</sc> regions were older and more evolved than the isolated masers. We argued that the original radio continuum survey was sensitive enough to detect UCH <sc>ii</sc> regions associated with the less luminous isolated masers (see Paper 1). However, if a sensitivity bias existed, it would manifest as a mass and luminosity offset between the radio loud and radio quiet, i.e. the UCH <sc>ii</sc> regions would be <italic>more massive and more luminous</italic> than isolated masers. In other words, some of the lower <italic>L</italic><sub>stars</sub> isolated masers may indeed be associated with UCH <sc>ii</sc> regions which fall below our detection limit. In Paper 1, we showed that most sources are luminous enough to have detectable UCH <sc>ii</sc> regions, assuming up to 90 per cent of the Lyman continuum photons are absorbed by dust (<xref ref-type="bibr" rid="b25">Kurtz, Churchwell & Wood 1994</xref>). If the disparity in luminosity were due to a real sensitivity limit, this would imply that, on average, at least 99 per cent of the ionizing photons are absorbed by dust before contributing to the creation of a UCH <sc>ii</sc> region. We believe this is unlikely in the majority of cases.</p>
<p>Some of the isolated maser sites may also be associated with Hyper-Compact H <sc>ii</sc> (HCH <sc>ii</sc>) regions. This recently discovered a class of H <sc>ii</sc> region is typically less than 0.01 pc across (<xref ref-type="bibr" rid="b2">Avalos et al. 2006</xref>) and exhibits a free–free emission spectrum with a turnover at higher frequencies than UCH <sc>ii</sc> regions, typically 20 GHz. HCH <sc>ii</sc> regions are expected to be optically thick at 8.6 GHz due to the squared dependence of the flux density with frequency and will be difficult to detect. <xref ref-type="bibr" rid="b28">Longmore et al. (2007)</xref> had used the Australia Telescope Compact Array to search for 22 GHz emission towards 24 of the objects presented here. Nine of these sources fall into the radio-quiet category and of these seven exhibit compact (30 → 3 arcsec) continuum emission. This result strongly suggests that some of our radio-quiet sources are associated with HCH <sc>ii</sc> regions, or indeed UCH <sc>ii</sc> regions undetected at 8.6 GHz. However, because HCH <sc>ii</sc> regions are thought to be the precursors of UCH <sc>ii</sc> regions, we are still selecting for two different evolutionary states.</p>
<p>The most obvious explanation for the offset is due to the time-scales on which we expect massive young stellar objects to progress to develop a detectable UCH <sc>ii</sc> region. In essence, clumps containing very high mass objects will evolve very rapidly and will spend only a short time in hot core and earlier phases. Because of the short time-scales involved, very massive hot cores will also likely be deeply embedded. The first indication of star formation activity may be when the expanding H <sc>ii</sc> region begins to disperse the surrounding molecular gas. Anecdotal evidence supporting this explanation is already emerging from uniform searches for massive young stellar objects (MYSOs), such as the Red MSX Source (RMS) survey (<xref ref-type="bibr" rid="b20">Hoare et al. 2004</xref>; <xref ref-type="bibr" rid="b33">Mottram 2008</xref>). The less evolved MYSOs identified in the survey have a significantly smaller upper mass cut-off compared to the UCH <sc>ii</sc> regions detected. No MYSO counterparts are found to the most massive UCH <sc>ii</sc> regions, implying that the UCH <sc>ii</sc> phase is the first indication of star formation in the most massive clumps.</p>
<p><italic>We believe differences between the radio-bright and radio-quiet samples may be attributed to evolution, with the caveat that the highest luminosity objects do not spend significant time in an ‘isolated maser’ phase.</italic></p>
</sec>
</sec>
<sec id="ss6">
<title>6 DETAILED COMPARISON BETWEEN SUB-SAMPLES</title>
<p>The aim of this molecular survey is to investigate the utility of molecular lines in the 3-mm band to act as ‘molecular clocks’. In this section, we consider the differences between subsamples divided from the source list based on associations with tracers of potentially different evolutionary states.</p>
<p>In Paper 1, we examined the link between hot cores, 6.67 GHz CH<sub>3</sub>OH masers and UCH <sc>ii</sc> regions. We found that CH<sub>3</sub>CN was brighter and more commonly detected towards masers associated with UCH <sc>ii</sc> regions, than towards ‘isolated’ masers. Asymmetries in the HCO<sup>+</sup> versus H<sup>13</sup>CO<sup>+</sup> line profiles were interpreted as inward (blue skewed) or outward (red skewed) gas motions. We also examined the associations with mid-infrared (mid-IR) emission in the Midcourse Space Experiment (MSX) images and found 11 dark clouds (seen in absorption against the 8 μm Galactic background) and 67 in emission. Many of the sources also exhibited high-velocity line wings, indicative of outflows. Based on these results, we separate the source list into subsamples based on the following criteria.</p>
<p><list id="l1" list-type="roman-lower">
<list-item><p>Presence or absence of radio emission from a UCH <sc>ii</sc> region.</p></list-item>
<list-item><p>CH<sub>3</sub>CN detected versus no CH<sub>3</sub>CN detected.</p></list-item>
<list-item><p>MSX dark (seen in absorption at 8 μm) versus MSX bright.</p></list-item>
<list-item><p>Presence or absence of high-velocity line wings in HCO<sup>+</sup>, HNC or <sup>13</sup>CO.</p></list-item>
<list-item><p>Presence or absence of blue-skewed HCO<sup>+</sup> line profiles, indicating infall.</p></list-item>
</list></p>
<p>In the following analysis, we contrast the subsamples by looking for differences in the median properties. A KS test is used to assess the similarity of the two distributions and trends are found based on large differences in the median values, coupled with a low probability that the subpopulations are drawn from the same parent distribution.</p>
<p>Initially, we compare the spectral line luminosities calculated from the integrated line intensity multiplied by the square of the distance to the source. If the distance is well known then comparisons of line luminosities will potentially reveal intrinsic differences between the subsamples. However, the error in the distance for this sample is relatively large (±0.65 kpc ≈ 20 per cent), and increases drastically towards Galactic longitudes <5°, where the model Galactic rotation curve begins to break down. Only the strongest trends will be seen using such a comparison.</p>
<p>We also compare the distance-independent linewidths. Linewidths in massive star-forming regions are generally much greater than the thermal width and are indicators of the level of turbulence present in the gas. Turbulence arises from bulk motions of gas, is often associated with energetic phenomena, such as outflows, and is correlated with star formation activity (e.g. <xref ref-type="bibr" rid="b14">Dobbs, Bonnell & Clark 2005</xref>).</p>
<p>The ratio of two spectral line intensities is unaffected by the distance and constitutes a much better probe of the physical and chemical conditions. The dominant errors are due to the relative accuracy with which the data have been calibrated and the unknown beam filling factors. The results presented here are based on data taken over four years and the calibration error in the data has been determined to be better than 30 per cent in most cases. The uncertainty due to the beam filling factors is given by the ratio of the solid angles over which the two molecules emit, (Ω<sub>1</sub>/Ω<sub>2</sub>), and may be large for dense-gas tracers which subtend angles smaller than the beam. The trends underlying the data should still be visible given the sample size.</p>
<p>Ideally, in a search for molecular clocks, the abundances should be compared directly. Unfortunately, the small number of transitions observed in this work makes it difficult to constrain the column densities accurately. Two or more transitions of a single molecule are needed to solve for the optical depth, excitation temperature and column density. We have been able to determine the column densities of thermal CH<sub>3</sub>CN and CH<sub>3</sub>OH through the rotational diagram method, and by assuming an excitation temperature or [<sup>12</sup>C]/[<sup>13</sup>C] abundance ratio, we have calculated the column densities of N<sub>2</sub>H<sup>+</sup> and HCO<sup>+</sup>. Here, we compare the beam-averaged column densities of these molecules directly to find the abundance ratios.</p>
<p>Lastly, we contrast the differences in a suite of measured physical parameters such as bolometric luminosity, gas mass, rotational temperature and mid-IR flux density, drawn from this and other work.</p>
<p>The important differences are summarized in <xref ref-type="sec" rid="ss6-4">Section 6.4</xref>.</p>
<sec id="ss6-1">
<title>6.1 Radio-loud versus radio-quiet sources</title>
<p><xref ref-type="fig" rid="f10">Fig. 10</xref> illustrates the difference in line luminosity between sources with and without 8 GHz radio continuum emission (from UCH <sc>ii</sc> regions) in the Mopra beam: the radio-loud and radio-quiet subsamples, respectively. As with CH<sub>3</sub>CN and HCO<sup>+</sup> in Paper 1, we find that the luminosity of the remaining transitions is enhanced towards the radio-loud subsample. The median-line luminosity is greater by 0.65σ, on average, where σ is the standard deviation of the combined sample. Thermal CH<sub>3</sub>OH stands out as having the most significantly different distribution and the largest difference in line luminosity, followed by H<sup>13</sup>CO<sup>+</sup>, CH<sub>3</sub>CN, HNC and N<sub>2</sub>H<sup>+</sup>.</p>
<fig position="float" id="f10">
<label>Figure 10</label>
<caption>
<p>The bars in this horizontal histogram represent the difference between the median values of two distributions, normalized to the standard deviation of the combined sample. Here, we compare <italic>spectral line luminosity</italic> in sources with and without UCH <sc>ii</sc> regions. A positive value means that the line is more luminous towards UCH <sc>ii</sc> regions. To assess the significance of the differences, we compute the KS statistic (left-hand column) and associated percentage probability (right-hand column). As a visual aid, the grey-scale of the bars is scaled to the probability. Probabilities smaller than approximately 2 per cent (dark bars) indicate significant differences.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f10.gif"></graphic>
</fig>
<p>We find that the median linewidths of all species are broader towards the radio-loud subsample (see <xref ref-type="fig" rid="f11">Fig. 11</xref>). CH<sub>3</sub>OH exhibits the greatest differences, closely followed by N<sub>2</sub>H<sup>+</sup> and H<sup>13</sup>CO<sup>+</sup>, all of which have statistically different distributions according to the KS test. HCO<sup>+</sup>, HCN, HNC and CH<sub>3</sub>CN have broader lines on average; however, the distributions are not significantly different. The differences in the median linewidth between the radio-loud and radio-quiet subsample range from 0.37 to 1.62 km s<sup>−1</sup>, greater than can be attributed to thermal broadening and most likely due to the greater turbulent energy in the gas associated with UCH <sc>ii</sc> regions.</p>
<fig position="float" id="f11">
<label>Figure 11</label>
<caption>
<p>Differences in the median linewidth, for sources with and without an associated UCH <sc>ii</sc> region. The linewidths were measured in two ways: by fitting a Gaussian to the line profile and measuring the FWHM or by integrating under the line and computing the equivalent width (HCN only). Other details are as in <xref ref-type="fig" rid="f10">Fig. 10</xref>.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f11.gif"></graphic>
</fig>
<p>The differences found between the linewidths and line luminosities are echoed in a comparison of the raw integrated line intensities. <xref ref-type="fig" rid="f12">Fig. 12</xref> shows the difference in the median values of the 28 intensity ratios between eight molecules. The median intensity of CH<sub>3</sub>OH and, to a lesser extent, CH<sub>3</sub>CN is clearly enhanced with respect to all other molecules. Statistically, the most significant difference in intensity ratio occurs between CH<sub>3</sub>OH and other molecules, specifically N<sub>2</sub>H<sup>+</sup> and <sup>13</sup>CO. The [CH<sub>3</sub>CN/<sup>13</sup>CO] and [H<sup>13</sup>CO<sup>+</sup>/<sup>13</sup>CO] ratios also stand out as being significantly greater towards the radio-loud subsample.</p>
<fig position="float" id="f12">
<label>Figure 12</label>
<caption>
<p>Differences in the median line–line intensity ratio for sources with and without an associated UCH <sc>ii</sc> region. For molecules with multiple transitions (e.g. CH<sub>3</sub>CN, CH<sub>3</sub>OH), we have summed the intensity of all measured lines.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f12.gif"></graphic>
</fig>
<p><xref ref-type="fig" rid="f13">Fig. 13</xref> illustrates a comparison of the median beam-averaged column densities of four molecules: CH<sub>3</sub>CN, CH<sub>3</sub>OH, HCO<sup>+</sup> and N<sub>2</sub>H<sup>+</sup>. Again, thermal CH<sub>3</sub>OH stands out as having a greater column density, and hence a higher abundance, towards the radio-loud subsample.</p>
<fig position="float" id="f13">
<label>Figure 13</label>
<caption>
<p>Differences in the median line–line column density for sources with and without an associated UCH <sc>ii</sc> region.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f13.gif"></graphic>
</fig>
<p>The median differences between other measured and calculated properties are shown in <xref ref-type="fig" rid="f14">Fig. 14</xref>. The median bolometric luminosity and median mass of the radio-loud subsample are significantly greater than for the isolated masers (see <xref ref-type="sec" rid="ss5">Section 5</xref>). The rotational temperature derived from CH<sub>3</sub>CN is similar in both groups, however, the CH<sub>3</sub>OH rotational temperature is significantly greater in the radio-bright group (also see <xref ref-type="fig" rid="f8">Fig. 8</xref>).</p>
<fig position="float" id="f14">
<label>Figure 14</label>
<caption>
<p>Differences in other computed and measured parameters for sources with and without an associated UCH <sc>ii</sc> region.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f14.gif"></graphic>
</fig>
</sec>
<sec id="ss6-2">
<title>6.2 Sources with and without CH<sub>3</sub>CN</title>
<p>CH<sub>3</sub>CN was detected towards 58 sources (70 per cent), however the detection rate towards the radio-loud subsample (19 sources) was 90 per cent. The enhanced luminosities and masses of our UCH <sc>ii</sc> regions may bias the results of a comparison between CH<sub>3</sub>CN-bright and CH<sub>3</sub>CN-dark groups. We therefore exclude any regions with radio emission from the comparison. In practise, we find that the major results are robust to the presence of UCH <sc>ii</sc> regions.</p>
<p>The differences in median line luminosity between the two groups presented in <xref ref-type="fig" rid="f15">Fig. 15</xref> are not as pronounced as in the previous section, however, all lines are marginally (0.07σ) more luminous towards the CH<sub>3</sub>CN-bright subsample. HCO<sup>+</sup> and HCN have the most significantly different distributions. We note that the HCO<sup>+</sup> and HCN profiles are affected by self-absorption, with the CH<sub>3</sub>CN-bright group containing the most extreme examples. The significant difference in the HCO<sup>+</sup> luminosity is also not reflected in the optically thin H<sup>13</sup>CO<sup>+</sup> isotopomer, suggesting that an abundance difference is unlikely. The greatest difference in the median values occurs in <sup>13</sup>CO, followed by CH<sub>3</sub>OH and H<sup>13</sup>CO<sup>+</sup>.</p>
<fig position="float" id="f15">
<label>Figure 15</label>
<caption>
<p>Differences in the median line luminosity for sources with and without CH<sub>3</sub>CN.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f15.gif"></graphic>
</fig>
<p>Linewidths are also biased towards being greater in the sample with detected CH<sub>3</sub>CN (see <xref ref-type="fig" rid="f16">Fig. 16</xref>). HCO<sup>+</sup> exhibits the most significant difference, followed by CH<sub>3</sub>OH, although a KS test cannot distinguish between the populations with confidence.</p>
<fig position="float" id="f16">
<label>Figure 16</label>
<caption>
<p>Differences in the median linewidth for sources with and without CH<sub>3</sub>CN.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f16.gif"></graphic>
</fig>
<p>When we examine the line intensity ratios in <xref ref-type="fig" rid="f17">Fig. 17</xref>, a similar picture emerges to the case of the UCH <sc>ii</sc> regions. Thermal CH<sub>3</sub>OH stands out, as its intensity relative to all other lines is enhanced in the CH<sub>3</sub>CN-bright subsample. We also see evidence for enhanced N<sub>2</sub>H<sup>+</sup> emission. In contrast, <sup>13</sup>CO emission is relatively weak towards the CH<sub>3</sub>CN-bright subsample. These ratios are highly suggestive of different environmental conditions in the dense gas (traced by N<sub>2</sub>H<sup>+</sup> and CH<sub>3</sub>OH), compared to the more extended gas (traced by <sup>13</sup>CO), and correlate well with the line luminosities.</p>
<fig position="float" id="f17">
<label>Figure 17</label>
<caption>
<p>Differences in the median line–line intensity ratio for sources with and without CH<sub>3</sub>CN. For line with multiple transitions, we summed the integrated intensity of all measured lines.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f17.gif"></graphic>
</fig>
<p><xref ref-type="fig" rid="f18">Fig. 18</xref> shows the differences in the median column densities of CH<sub>3</sub>OH, HCO<sup>+</sup> and N<sub>2</sub>H<sup>+</sup>. Again, the thermal CH<sub>3</sub>OH column is seen to be significantly enhanced towards the CH<sub>3</sub>CN bright group, echoing the result from comparing the raw-line ratios.</p>
<fig position="float" id="f18">
<label>Figure 18</label>
<caption>
<p>Differences in the median column densities for sources with and without CH<sub>3</sub>CN.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f18.gif"></graphic>
</fig>
<p>The difference in the median values of other derived parameters is displayed in <xref ref-type="fig" rid="f20">Fig. 20</xref>. Sources with detected CH<sub>3</sub>CN have significantly brighter thermal continuum emission at 1.2-mm wavelengths. This translates into a mean mass higher by ∼300 M<sub>⊙</sub>, omitting UCH <sc>ii</sc> regions or ∼600 M<sub>⊙</sub>, including UCH <sc>ii</sc> regions. Histograms of the mass distributions for the two groups are plotted in <xref ref-type="fig" rid="f19">Fig. 19</xref> (left-hand panel), and show significant differences. All other parameters presented in <xref ref-type="fig" rid="f20">Fig. 20</xref> are similar for both subsamples.</p>
<fig position="float" id="f20">
<label>Figure 20</label>
<caption>
<p>Differences in other computed and measured parameters for sources with and without CH<sub>3</sub>CN.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f20.gif"></graphic>
</fig>
<fig position="float" id="f19">
<label>Figure 19</label>
<caption>
<p>Left-hand panel: histograms showing the distributions of gas masses derived from 1.2-mm continuum flux density for sources with (top panel) and without (bottom panel) detected CH<sub>3</sub>CN. Sources in which CH<sub>3</sub>CN has been detected tend to be more massive. Right-hand panel: the distribution of distances for sources with (top panel) and without (bottom panel) detected CH<sub>3</sub>CN. A tail of more distant objects exists in the CH<sub>3</sub>CN-dark sample, suggesting that a small fraction of non-detections may be due to sensitivity limits.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f19.gif"></graphic>
</fig>
<p>The division of the sample into populations with and without detected CH<sub>3</sub>CN may be a somewhat artificial distinction. While CH<sub>3</sub>CN emission may be present in more distant sources, it may fall below our 2σ detection threshold of 80 mK. Emission over a small solid angle will suffer from high beam dilution factors, further compounding the problem. The distributions of distance in the two samples, plotted in <xref ref-type="fig" rid="f19">Fig. 19</xref> (right-hand panel), are similar, except for a tail of more distant objects without CH<sub>3</sub>CN. This difference suggests that a small fraction of non-detections may be attributable to sensitivity limits.</p>
<p>We also split the sample via other criteria (see the Introduction to this section), however, the chemical and physical differences are not as significant. The results of these investigations are presented as an online supplement.</p>
</sec>
<sec id="ss6-3">
<title>6.3 The effect of beam dilution</title>
<p>As mentioned in <xref ref-type="sec" rid="ss4-6">Section 4.6</xref>, the beam filling factors constitute an unknown source of error when comparing the intensities or column densities of two molecules. If both emitting regions subtend angles (θ) smaller than the beam, then we must correct for the relative volume filling factor given by (θ<sub>1</sub>/θ<sub>2</sub>)<sup>2</sup>. Assuming the emitting regions have similar characteristic sizes in all sources, the relative beam filling factors will average to ∼1 over the entire sample. However, if one species generally fills the beam while the other species is generally beam-diluted (i.e. one has constant angular size while the other does not, then we would expect a 1 distance<sup>−2</sup> dependence in the line intensity ratios). <xref ref-type="fig" rid="f21">Fig. 21</xref> presents two plots of line intensity ratio versus distance. The left-hand panel plots the CH<sub>3</sub>CN/CH<sub>3</sub>OH ratio versus distance. The emission from both molecules is expected to derive from regions smaller than the 36 arcsec Mopra beam, even for nearby sources. We see the line ratios are not correlated with distance, consistent with a constant ratio between the angular sizes of the species. The right-hand panel plots the CH<sub>3</sub>CN/N<sub>2</sub>H<sup>+</sup> ratio versus distance. Assuming N<sub>2</sub>H<sup>+</sup> emission closely follows the dust morphology (e.g. <xref ref-type="bibr" rid="b7">Caselli et al. 2002a</xref>; <xref ref-type="bibr" rid="b36">Pirogov et al. 2003</xref>), we have shown in Table S9 that the beam dilution factor is generally close to 1. In <xref ref-type="fig" rid="f21">Fig. 21</xref> (right-hand panel), the dashed line plots the expected effect of CH<sub>3</sub>CN beam dilution as a source is moved further away. We see the scatter in measured line ratio increases with distance making it difficult to interpret the data and, in addition to beam dilution, chemical and physical differences in the sources affect the line intensity ratios. Omitting the UCH <sc>ii</sc> regions (squares) and considering only isolated masers (crosses), we see a possible trend towards decreasing CH<sub>3</sub>CN/N<sub>2</sub>H<sup>+</sup> ratio with increasing distance. This trend does not approach the 1 distance<sup>−2</sup> dependence expected, but it may indicate that some of our comparisons will include beam dilution effects for CH<sub>3</sub>CN.</p>
<fig position="float" id="f21">
<label>Figure 21</label>
<caption>
<p>Plots of line–line intensity ratio versus distance. (left-hand panel) [CH<sub>3</sub>CN/CH<sub>3</sub>OH] ratio versus distance. (right-hand panel) [CH<sub>3</sub>CN/N<sub>2</sub>H<sup>+</sup>] ratio versus distance. Crosses mark isolated masers, while squares mark UCH <sc>ii</sc> regions. The dotted line shows how the [CH<sub>3</sub>CN/N<sub>2</sub>H<sup>+</sup>] ratio is expected to decrease with distance if the CH<sub>3</sub>CN emission has a fixed physical size and the N<sub>2</sub>H<sup>+</sup> emission always fills the beam.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f21.gif"></graphic>
</fig>
</sec>
<sec id="ss6-4">
<title>6.4 Interpretation and discussion</title>
<p>It is clear we are seeing significant differences between sources associated with different tracers. In this section, we summarize the differences between the subsamples and suggest some possible interpretations of the results.</p>
<sec id="ss6-4-1">
<title>6.4.1 UCH <sc>ii</sc> regions versus isolated masers</title>
<p>The clearest distinctions exist between isolated CH<sub>3</sub>OH masers and those with UCH <sc>ii</sc> regions. The result found in Paper 1, where the CH<sub>3</sub>CN and HCO<sup>+</sup> lines were shown to be more luminous and have greater linewidths in the presence of a UCH <sc>ii</sc> region, has been extended to the other observed species. When we examine the raw line intensity ratios, thermal CH<sub>3</sub>OH distinguishes itself as being relatively intense towards the radio-loud sources. This is reflected in both the higher median column density and CH<sub>3</sub>OH rotational temperature. We interpret this result as evidence for an enhanced abundance of CH<sub>3</sub>OH in relatively warm gas nearby the UCH <sc>ii</sc> regions. Significant differences are also found between other tracers of high- and low-density gas. The [CH<sub>3</sub>CN/<sup>13</sup>CO], [N<sub>2</sub>H<sup>+</sup>/<sup>13</sup>CO] and [H<sup>13</sup>CO<sup>+</sup>/<sup>13</sup>CO] ratios are enhanced in the radio-loud subsample, possibly indicating different excitation conditions or abundances in the ‘core’ compared to the ‘envelope’ gas. Interestingly, when compared to all other lines, the relative brightness of the hot-core tracer CH<sub>3</sub>CN is <italic>not</italic> significantly greater, yet we detect it towards 95 per cent of UCH <sc>ii</sc> regions and only 63 per cent of the isolated masers.</p>
<p>A potential explanation for the median offset in line luminosity and linewidths between the radio-quiet and radio-loud samples may be found in the significantly greater bolometric luminosity and gas mass of the radio-loud sample. Simply put, more luminous clumps will likely have a greater flux of IR-photons and higher gas temperatures. Under these conditions, one would expect to see more luminous spectral lines. Linewidth is commonly used as a proxy for star formation activity, and hence evolutionary state, as it is largely dependent on turbulence in massive star-forming regions. However, on the scale of molecular clouds (3–20 pc), the linewidth has also been empirically correlated with the cloud mass (<xref ref-type="bibr" rid="b43">Solomon et al. 1987</xref>), possibly due to multiple (sub)clumps with different velocities within the beam. It is unclear which effect dominates in our sample. <xref ref-type="fig" rid="f22">Fig. 22</xref> plots the linewidth N<sub>2</sub>H<sup>+</sup> versus the gas mass. N<sub>2</sub>H<sup>+</sup> is chosen as it is optically thin and has been shown to trace the dense regions of the molecular clouds (e.g. <xref ref-type="bibr" rid="b8">Caselli et al. 2002b</xref>). A weak correlation is observed which may account for some of the linewidth differences between the groups, however we cannot rule out an evolutionary effect.</p>
<fig position="float" id="f22">
<label>Figure 22</label>
<caption>
<p>Plot of N<sub>2</sub>H<sup>+</sup> FWHM linewidth versus gas mass. The squares represent UCH <sc>ii</sc> regions and crosses represent isolated masers.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f22.gif"></graphic>
</fig>
<p>That the relative intensity and abundance of CH<sub>3</sub>OH is greater towards the radio-loud sources likely reflects real chemical differences between the subsamples. This may be interpreted as evidence for the more advanced evolutionary state of the radio-loud sample. The CH<sub>3</sub>OH abundance is predicted to be enhanced by a number of mechanisms as a massive young stellar object evolves. In the hot-core phase, it is known to be evaporated from the dust grains as the heating source evolves (<xref ref-type="bibr" rid="b11">Charnley et al. 1995</xref>). The abundance of CH<sub>3</sub>OH is also predicted to be enhanced in the walls of outflow cavities, where it is sputtered from the grains by low-velocity shocks (<xref ref-type="bibr" rid="b21">Hogerheijde 2005</xref>).</p>
</sec>
<sec id="ss6-4-2">
<title>6.4.2 CH<sub>3</sub>CN bright versus CH<sub>3</sub>CN dark</title>
<p>In this and the remaining comparisons, we removed the radio-loud sources from the list, as their higher intrinsic luminosities may bias the results. We find sources with detected CH<sub>3</sub>CN are also biased towards having greater line luminosities and linewidths, although these differences are not as pronounced as with the radio-loud versus radio-quiet division. The most striking result is the significant tendency for thermal CH<sub>3</sub>OH to be brighter and have a greater column density in the CH<sub>3</sub>CN bright group. This is not unexpected as the CH<sub>3</sub>CN and CH<sub>3</sub>OH abundances are predicted to increase with time in hot cores (<xref ref-type="bibr" rid="b11">Charnley et al. 1995</xref>). We also see evidence for more turbulent and energetic conditions in the dense gas (traced by N<sub>2</sub>H<sup>+</sup> and CH<sub>3</sub>OH) compared to more extended and diffuse gas (traced by <sup>13</sup>CO), implying a more advanced evolutionary state in the CH<sub>3</sub>CN-bright subsample. Interestingly, the CH<sub>3</sub>CN-bright subsample also has a significantly greater median mass, which may bias this result. CH<sub>3</sub>CN is also not detected in sources with distances greater than 10 kpc, suggesting that a small number of sources may contain CH<sub>3</sub>CN below the sensitivity limit of our survey.</p>
</sec>
<sec id="ss6-4-3">
<title>6.4.3 MSX bright versus MSX dark</title>
<p>All species in the MSX-dark sample have marginally greater line luminosities and linewidths compared to their MSX-bright counterparts, with the <sup>13</sup>CO linewidth being especially notable. This is unexpected, as IR dark clouds have been postulated to be the young and quiescent precursors to hot cores and UCH <sc>ii</sc> regions (<xref ref-type="bibr" rid="b38">Rathborne, Jackson & Simon 2006</xref>). HCO<sup>+</sup> and CH<sub>3</sub>OH tend to be brighter and CH<sub>3</sub>OH marginally more abundant in the MSX-dark subsample, although this is not judged to be significant. Other parameters are similar in both subsamples. Given the small number of sources in the MSX-bright subsample (9–11 sources), none of the observed differences is large enough to draw any definite conclusions from the data.</p>
</sec>
<sec id="ss6-4-4">
<title>6.4.4 Line wings versus no line wings</title>
<p>There are no significant differences between the median line luminosities when the sources are divided into subsamples with and without high-velocity line wings. Differences between the HCO<sup>+</sup> linewidth distributions can be attributed to unidentified low-velocity line wings. <sup>13</sup>CO is significantly brighter (relative to other lines) in the sources without line wings and CH<sub>3</sub>CN is biased towards being brighter and more abundant in the subsample with line wings. No distinguishing differences are seen in the physical parameters or IR-colour ratios.</p>
</sec>
<sec id="ss6-4-5">
<title>6.4.5 Infall profiles versus no infall profiles</title>
<p>The properties of sources with and without HCO<sup>+</sup>‘infall profiles’ are close to identical. Little difference is seen between the median line luminosities, linewidths or line intensity ratios. Biases in the H<sup>13</sup>CO<sup>+</sup> and HCO<sup>+</sup> line ratios can be attributed to selection effects.</p>
<p>The online material contains additional discussion and figures pertaining to the previous three paragraphs.</p>
</sec>
</sec>
<sec id="ss6-5">
<title>6.5 Further discussion</title>
<p>Is the mass of the hot core related to the dust mass?<xref ref-type="fig" rid="f23">Fig. 23</xref> shows a plot of CH<sub>3</sub>CN beam-averaged column density versus mass of gas derived from 1.2-mm continuum emission. A clear positive correlation is seen between the two. If we assume the column of CH<sub>3</sub>CN is a proxy for the mass of the hot core, then it follows that the mass of the hot core increases with the mass of the extended clump. The relationship is well fit by <inline-formula><inline-graphic xlink:href="mnras0394-0323-mu10.gif"></inline-graphic></inline-formula> and assuming the density is constant within the gas, this implies that <italic>M</italic><sub>hot core</sub>∝<italic>M</italic><sup>0.5</sup><sub>dust</sub>. This is consistent with the results of <xref ref-type="bibr" rid="b9">Cesaroni (2005)</xref> who found that the hot core is an continuation of the clump, and there is no break in the approximately 1/<italic>r</italic><sup>2</sup> density power law. Rather, the hot core is differentiated by temperature and excitation conditions influenced by the central heating source.</p>
<fig position="float" id="f23">
<label>Figure 23</label>
<caption>
<p>Plot of CH<sub>3</sub>CN column density versus gas mass derived from 1.2-mm continuum flux density.</p>
</caption>
<graphic xlink:href="mnras0394-0323-f23.gif"></graphic>
</fig>
</sec>
</sec>
<sec id="ss7">
<title>7 SUMMARY AND CONCLUSIONS</title>
<p>In continuation of the ‘Hot Molecular Cores’ survey, we used the Mopra telescope to search for 3-mm transitions of CH<sub>3</sub>OH (2–1), N<sub>2</sub>H<sup>+</sup> (1–0), <sup>13</sup>CO (1–0), HCN (1–0) and HNC (1,– 0). This is in addition to the previously reported measurements of HCO<sup>+</sup> (1–0), H<sup>13</sup>CO<sup>+</sup> (1–0) and CH<sub>3</sub>CN (5–4) and (6–5), in Paper 1. Molecular emission was detected in all but one source (G10.10−0.72), with N<sub>2</sub>H<sup>+</sup>, <sup>13</sup>CO, HNC and HCN detected towards all of the remaining 82 targets. Thermal CH<sub>3</sub>OH emission was detected in 78 sources (94 per cent). The following conclusions have been drawn from the data.</p>
<p><list id="l2" list-type="roman-lower">
<list-item><p>Virial masses estimated from N<sub>2</sub>H<sup>+</sup> are comparable to or lower than the gas masses of the sample. Values for <italic>M</italic><sub>gas</sub> have been calculated from the 1.2-mm dust emission and range from 30 to 5 × 10<sup>3</sup> M<sub>⊙</sub>, while values for <italic>M</italic><sub>virial</sub> range from ∼40 to 2.5 × 10<sup>3</sup> M<sub>⊙</sub>. We note that N<sub>2</sub>H<sup>+</sup> exhibits particularly narrow linewidths with an average of 3.0 km s<sup>−1</sup>.</p></list-item>
<list-item><p>H<sub>2</sub> volume densities have also been estimated from the surface brightness and extent of 1.2-mm emission. The mean value is 5 × 10<sup>4</sup> cm<sup>−3</sup>, however, we note that the assumptions inherent in the calculation introduce large unknown uncertainties in the values for individual sources.</p></list-item>
<list-item><p>Rotational temperatures derived from CH<sub>3</sub>OH range from 3.0 to 14.0 K, with a mean value of 6.67 K. These anomalously low temperatures strongly suggest that the CH<sub>3</sub>OH is subthermally excited. Assuming a common excitation temperature, we find that the abundance ratio of A- to E-type CH<sub>3</sub>OH ranges from 0.5 to 2.8, with a mean of 1.5, consistent with previously published values.</p></list-item>
<list-item><p>N<sub>2</sub>H<sup>+</sup> is found to be optically thin towards most of the sample. Likewise the CH<sub>3</sub>OH rotational diagrams are consistent with low optical depths in the <italic>J</italic> = 2–1 transitions. <sup>13</sup>CO, HCN and HNC all exhibit line profile asymmetries and may be significantly optically thick.</p></list-item>
<list-item><p>The <italic>luminosity/gas–mass</italic> relationship for the sample is found to follow a rough power law of the form <italic>L</italic>∝<italic>M</italic><sup>0.68</sup>, meaning that the lower mass clumps are overluminous and the higher mass clumps are underluminous. We find that the gas mass of the clumps is comparable to or greater than the mass of the stellar content (implied by the luminosity). We interpret this as an indicator of young ages, as the newborn massive stars still retain their natal molecular clouds.</p></list-item>
<list-item><p>The radio-loud subsample (UCH <sc>ii</sc> regions) is on average more massive <italic>and</italic> more luminous than the radio-quiet subsample (isolated masers). We argue that the most likely explanation for this disparity is due to the comparatively rapid time-scale over which the most massive objects evolve into H <sc>ii</sc> regions. The first significant tracer of star formation in these objects may be radio-continuum emission from a HCH <sc>ii</sc> or UCH <sc>ii</sc> region. It is also possible that some of the isolated maser sample contains HCH <sc>ii</sc> regions which are too optically thick to detect at 8 GHz. In either case, there is a clear age difference between the subsamples.</p></list-item>
<list-item><p>We separated the source list into subsamples associated with different tracers of evolution and compared their median physical and chemical properties. Our findings are as follows.</p>
<list list-type="alpha-lower">
<list-item><p><italic>All</italic> spectral lines are brighter and more luminous towards the radio-loud sample, extending the result found in Paper 1.</p></list-item>
<list-item><p>The thermal CH<sub>3</sub>OH brightness and abundance is found to be enhanced (compared to other species) in radio-loud and CH<sub>3</sub>CN-bright sources. We suggest that in single-dish observations, thermal CH<sub>3</sub>OH may be used as a crude molecular clock to indicate the evolutionary state of a clump.</p></list-item>
<list-item><p>Given the small sample size, conclusions drawn from a comparison of MSX-bright and MSX-dark sources are not significant. It is worth noting, however, that linewidths and line luminosities are generally greater in the IR dark clouds, indicating perhaps that when selected using CH<sub>3</sub>OH masers, dark clouds conceal objects at later stages of evolution.</p></list-item>
<list-item><p>We divided the source list into subsamples with and without high-velocity line wings and HCO<sup>+</sup>-infall profiles, but found no significant differences in chemical or physical properties.</p></list-item>
</list></list-item>
<list-item><p>We find that the beam-averaged CH<sub>3</sub>CN column density is correlated with the gas mass derived from 1.2-mm emission. This suggests that the mass of the hot core is dependent on the mass of the clump in which it is embedded.</p></list-item>
</list></p>
</sec>
</body>
<back>
<fn-group>
<fn id="fn1">
<label>1</label><p><sc>spc</sc> is a spectral line reduction package for single dish telescopes written by the ATNF. <ext-link ext-link-type="uri" xlink:href="http://www.atnf.csiro.au/software/">http://www.atnf.csiro.au/software/</ext-link></p>
</fn>
<fn id="fn2">
<label>2</label><p><sc>xs</sc> is written by P. Bergman, Onsala Space Observatory. <ext-link ext-link-type="uri" xlink:href="ftp://yggdrasil.oso.chalmers.se/pub/xs/">ftp://yggdrasil.oso.chalmers.se/pub/xs/</ext-link></p>
</fn>
<fn id="fn3">
<label>3</label><p><sc>class</sc> is part of the Grenoble Image and Line Data Analysis Software (GILDAS) working group software. <ext-link ext-link-type="uri" xlink:href="http://www.iram.fr/IRAMFR/GILDAS/">http://www.iram.fr/IRAMFR/GILDAS/</ext-link></p>
</fn>
<fn id="fn4">
<label>4</label><p>We have the following sources in common with Cesaroni: G5.89−0.39, G9.62+0.19, G10.47+0.03, G10.62−0.38, G29.96−0.02 and G31.41+0.31.</p>
</fn>
</fn-group>
<ack>
<p>The authors are grateful to the Australian Research Council and UNSW for grant support. CRP and SNL were supported by a University of New South Wales School of Physics Scholarship during the course of this research. The Mopra radio telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operations as a National Facility managed by CSIRO. During 2002-2005, the Mopra telescope was operated through a collaborative arrangement between the University of New South Wales and the CSIRO.</p>
<p>This research has made use of NASA's Astrophysics Data System.</p>
<p>We would also like to thank the anonymous referee for being especially thorough in seeking out mistakes in the paper. His/her comments were extremely useful and have helped to improve the presentation and interpretation.</p>
</ack>
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<title>Supporting Information</title>
<p><bold>Appendix A.</bold> Maser-selected hot cores and UCHii Regions.</p>
<p>Please note: Wiley-Blackwell is not responsible for the content or functionality of any supporting information supplied by the authors. Any queries (other than missing material) should be directed to the corresponding author for the article.</p>
</sec>
</back>
</article></istex:document></istex:metadataXml><mods version="3.6"><titleInfo><title>Physical and chemical conditions in methanol maser selected hot cores and UCH ii regions</title></titleInfo><titleInfo type="alternative" contentType="CDATA"><title>Physical and chemical conditions in methanol maser selected hot cores and UCH ii regions</title></titleInfo><name type="personal"><namePart type="given">C. R.</namePart><namePart type="family">Purcell</namePart><affiliation>University of Manchester, Jodrell Bank Observatory, Macclesfield, Cheshire SK11 9DL</affiliation><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>E-mail: Cormac.Purcell@manchester.ac.uk</affiliation><affiliation></affiliation><affiliation>E-mail: Cormac.Purcell@manchester.ac.uk</affiliation><role><roleTerm type="text">author</roleTerm></role></name><name type="personal"><namePart type="given">S. N.</namePart><namePart type="family">Longmore</namePart><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Harvard–Smithsonian Center For Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA</affiliation><affiliation>CSIRO Australia Telescope National Facility, PO Box 76, Epping, NSW 1710, Australia</affiliation><role><roleTerm type="text">author</roleTerm></role></name><name type="personal"><namePart type="given">M. G.</namePart><namePart type="family">Burton</namePart><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><role><roleTerm type="text">author</roleTerm></role></name><name type="personal"><namePart type="given">A. J.</namePart><namePart type="family">Walsh</namePart><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Centre for Astronomy, James Cook University, Townsville, QLD 4811, Australia</affiliation><role><roleTerm type="text">author</roleTerm></role></name><name type="personal"><namePart type="given">V.</namePart><namePart type="family">Minier</namePart><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Service d'Astrophysique, DAPNIA/DSM/CEA Saclay, 91191 Gif-sur-Yvette, France</affiliation><affiliation>AIM, Unité Mixte de Recherche, CEA-CNRS-Université Paris VII, UMR 7158, CEA/Saclay, 91191 Gif-sur-Yvette, France</affiliation><role><roleTerm type="text">author</roleTerm></role></name><name type="personal"><namePart type="given">M. R.</namePart><namePart type="family">Cunningham</namePart><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><role><roleTerm type="text">author</roleTerm></role></name><name type="personal"><namePart type="given">R.</namePart><namePart type="family">Balasubramanyam</namePart><affiliation>School of Physics, University of New South Wales, Sydney, NSW 2052, Australia</affiliation><affiliation>Raman Research Institute, Sadashivanagar, Bangalore 560 080, India</affiliation><role><roleTerm type="text">author</roleTerm></role></name><typeOfResource>text</typeOfResource><genre type="research-article" displayLabel="research-article" authority="ISTEX" authorityURI="https://content-type.data.istex.fr" valueURI="https://content-type.data.istex.fr/ark:/67375/XTP-1JC4F85T-7">research-article</genre><originInfo><publisher>Blackwell Publishing Ltd</publisher><place><placeTerm type="text">Oxford, UK</placeTerm></place><dateIssued encoding="w3cdtf">2009-03-21</dateIssued><dateCreated encoding="w3cdtf">2009-03-07</dateCreated><copyrightDate encoding="w3cdtf">2009</copyrightDate></originInfo><abstract>We present the results of a targeted 3-mm spectral line survey towards the eighty-three 6.67 GHz methanol maser selected star-forming clumps observed by Purcell. In addition to the previously reported measurements of HCO+ (1–0), H13CO+ (1–0) and CH3CN (5–4) and (6–5), we used the Mopra antenna to detect emission lines of N2H+ (1–0), HCN (1–0) and HNC (1–0) towards 82/83 clumps (99 per cent), and CH3OH (2–1) towards 78/83 clumps (94 per cent). The molecular line data have been used to derive virial and local thermodynamic equilibrium masses, rotational temperatures and chemical abundances in the clumps, and these properties have been compared between subsamples associated with different indicators of evolution. The greatest differences are found between clumps associated with 8.6 GHz radio emission, indicating the presence of an Ultra-Compact H ii (UCH ii) region, and ‘isolated’ masers (without associated radio emission), and between clumps exhibiting CH3CN emission and those without. In particular, thermal CH3OH is found to be brighter and more abundant in UCH ii regions and in sources with detected CH3CN, and may constitute a crude molecular clock in single dish observations. Clumps associated with 8.6 GHz radio emission tend to be more massive and more luminous than clumps without radio emission. This is likely because the most massive clumps evolve so rapidly that a Hyper-Compact H ii or UCH ii region is the first visible tracer of star formation. The gas mass to submm/infrared luminosity relation for the combined sample was found to be L∝M0.68, considerably shallower than expected for massive main-sequence stars. This implies that the mass of the clumps is comparable to, or greater than, the mass of the stellar content. We also find that the mass of the hot core is correlated with the mass of the clump in which it is embedded.</abstract><subject><genre>keywords</genre><topic>surveys</topic><topic>stars: formation</topic><topic>stars: pre-main-sequence</topic><topic>ISM: abundances</topic><topic>ISM: molecules</topic></subject><relatedItem type="host"><titleInfo><title>Monthly Notices of the Royal Astronomical Society</title></titleInfo><titleInfo type="abbreviated"><title>Monthly Notices of the Royal Astronomical Society</title></titleInfo><genre type="journal" authority="ISTEX" authorityURI="https://publication-type.data.istex.fr" valueURI="https://publication-type.data.istex.fr/ark:/67375/JMC-0GLKJH51-B">journal</genre><subject><topic>Papers</topic></subject><identifier type="ISSN">0035-8711</identifier><identifier type="eISSN">1365-2966</identifier><identifier type="PublisherID">mnras</identifier><identifier type="PublisherID-hwp">mnras</identifier><part><date>2009</date><detail type="volume"><caption>vol.</caption><number>394</number></detail><detail type="issue"><caption>no.</caption><number>1</number></detail><extent unit="pages"><start>323</start><end>339</end></extent></part></relatedItem><identifier type="istex">D9AEBE63A9901DB9EBE267E3263BC9DC7B158841</identifier><identifier type="DOI">10.1111/j.1365-2966.2008.14283.x</identifier><accessCondition type="use and reproduction" contentType="copyright">© 2009 The Authors. 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