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A&A
Volume 609, January 2018
Article Number A129
Number of page(s) 19
Section Interstellar and circumstellar matter
DOI https://doi.org/10.1051/0004-6361/201730576
Published online 01 February 2018

© ESO, 2018

1. Introduction

Nitrogen is the fifth most abundant element in the Universe. It possesses two stable isotopes, 14N and 15N. Füri & Marty (2015) proposed that there are three distinct isotopic reservoirs in the solar system: the protosolar nebula (PSN), the inner solar system, and cometary ices. In the terrestrial atmosphere (TA), the typical isotopic ratio 14N/15N as derived from N2 is ~272 (Marty et al. 2009). This value is almost a factor two higher than the ratio measured in nitrile-bearing molecules and in nitrogen hydrides of some comets (e.g. ~150, Manfroid et al. 2009; Shinnaka et al. 2016), and higher even by a factor five in carbonaceus chondrites (e.g. ≳50, Bonal et al. 2009). On the other hand, the ratio measured in solar wind particles collected by the Genesis spacecraft, representative of the PSN value and similar to the ratio measured on Jupiter (Owen et al. 2001), is 441±6 (Marty et al. 2010). Thus, the PSN 14N/15N value is about twice higher than the TA value and three times higher than the value measured in comets. These measurements suggest that multiple isotopic reservoirs were present very early in the formation process of the solar system. This was recently demonstrated based on the CN/C15N and HCN/HC15N isotopic ratios measured in a sample of PSN analogues (Hily-Blant et al. 2017). The nature and origin of these reservoirs remain elusive, however. Currently, the two main possibilities are i) isotope-selective photodissociation of N2 in the PSN by stellar UV (Heays et al. 2014; Guzman et al. 2017); and ii) an interstellar origin by mass fractionation reaction (Hily-Blant et al. 2013b; Roueff et al. 2015). A similar isotopic enrichment in comets, some meteorites, and interplanetary dust particles (IDPs) with respect to the PSN value has also been found for hydrogen. In the PSN, the D/H ratio is ~10-5 (Geiss et al. 1998), that is, similar to the cosmic elemental abundance (Linsky et al. 2006). In comets, different values of the D/H ratio were estimated: in the Jupiter-family comet 67P/Churyumov-Gerasimenko, the D/H ratio measured in H2O is about 5×10-4, approximately three times that of Earth’s oceans (Altwegg et al. 2015), and other results from Herschel have shown D/H ~ 1.5×10-4 in another Jupiter-family comet, 103P/Hartley (Hartogh et al. 2011), which is same as the ratio of Earth’s ocean. In carbonaceous chondrites, values of D/H ~ 1.2–2.2×10-4 in hydrous silicates were obtained (Robert 2003). Furthermore, very high D/H ratios of ~10-2 (Remusat et al. 2009) have been found in small regions in insoluble organic matter (IOM) of meteorites and IDPs; these regions are called “hot spots”. Deuterium fractionation in the interstellar medium and in solar system objects was also discussed by Ceccarelli et al. (2014).

From a theoretical point of view, the D/H enhancement for several species originates in a low-temperature environment where ion-molecule reactions are favoured, starting from the reaction

H3++HD→H2D++H2+232K,$$ {\rm H}_{3}^{+}+{\rm HD} \rightarrow {\rm H}_{2}{\rm D}^{+}+{\rm H}_{2}+{\rm 232~K} , $$which produces an enhanced H2D+/H+3\hbox{$_{3}^{+}$} abundance ratio for temperatures lower than ~50 K and when H2 is mainly in para form (e.g. Pagani et al. 1992; Gerlich et al 2002; Walmsley et al. 2004). Moreover, in dense cores, where CO freezes out onto dust grains (e.g. Caselli et al. 1999; Fontani et al. 2012), H2D+ can survive and H2D+/H+3\hbox{$_{3}^{+}$} is further enhanced causing a high level of deuteration in molecules. The role of grain surface chemistry on icy material during the early cold phase is also expected to play an important role for the deuteration of neutral species such as water, formaldehyde, methanol, and complex organic molecules (e.g. Cazaux et al. 2011; Taquet et al. 2012, 2013). The reason is that the enhanced abundances of the deuterated forms of H+3\hbox{$_{3}^{+}$} (H2D+, D2H+, D+3\hbox{$_{3}^{+}$}) produce enhanced abundances of D atoms in the gas phase upon their dissociative recombination with electrons. The enhanced D abundance implies a higher D/H ratio, so that unsaturated molecules on the surface (in particular CO) can be deuterated as well as hydrogenated to produce singly and doubly deuterated water and formaldehyde as well as singly, doubly, and triply deuterated methanol (e.g. Caselli & Ceccarelli 2012; Ceccarelli et al. 2014).

In principle, similar gas-phase mass fractionation reactions can also produce 15N enrichments under cold and dense conditions. This motivated the search for correlated enrichments in D and 15N in cosmomaterials (Aléon et al. 2009), although chemical models suggest that such correlations may not be present (Wirström et al. 2012). Furthermore, while the above scenario of D-enrichments through ion-neutral mass fractionation reactions has been firmly proved for D through observations of low- and high-mass star-forming regions (e.g. Crapsi et al. 2005; Caselli et al. 2008; Emprechtinger et al. 2009; Fontani et al. 2011), the reasons for 15N enrichment are still highly uncertain. For example, HC15N and H15NC could be formed through the dissociative recombination of HC15NH+: Terzieva & Herbst (2000) found that the reaction that causes most of the N-fractionation is the exchange reaction between 15N and HCNH+,

15N+HCNH+→N+HC15NH++35.1K,$$ ^{15}{\rm N} + {\rm HCNH}^{+} \rightarrow {\rm N} + {\rm HC}^{15}{\rm NH}^{+} + 35.1{\rm K}, $$but they assumed that this reaction could occur without an energy barrier. The most recent and complete chemical models (Roueff et al. 2015) are implemented with the recent discovery that this reaction has an energy barrier, and they indicate that 15N should not be enriched during the evolution of a star-forming core, not even at the very early cold phases. Moreover, observational works devoted to testing these predictions are still very limited.

In the few low-mass pre-stellar cores or protostellar envelopes observed so far, values of 14N/15N comparable to the value of the PSN have been measured from N2H+ and NH3 (330±150 in N2H+, Daniel et al. 2016; 350–500, Gerin et al. 2009; 334±50, Lis et al. 2010 in NH3), or even higher (1000±200 in N2H+, Bizzocchi et al. 2013). Conversely, through observations of the nitrile-bearing species HCN and HNC in low-mass sources, the 14N/15N ratio was found to be significantly lower (140–360, Hily-Blant et al. 2013a; 160–290, Wampfler et al. 2014). Moreover, low values have recently also been found in protoplanetary discs (from 80 up to 160, Guzmán et al. 2017), but these results are based on a low statistics (six protoplanetary discs). From the point of view of chemical models, differential fractionation is expected between nitriles and hydrides (Wirström et al. 2012; Hily-Blant et al. 2013a,b), although the most recent models (Roueff et al. 2015) do not support this scenario. However, none of these models is able to reproduce the low fractionation observed in N2H+ towards L1544 (Bizzocchi et al. 2013) and the large spread reported by Fontani et al. (2015a).

To investigate the possible correlation between D and 15N enrichments in interstellar environments, Fontani et al. (2015a) conducted a survey of the 14N/15N and D/H isotopic ratios towards a sample of high-mass cores in different evolutionary stages and found a possible anti-correlation, which was quite faint, however, and was obtained from a low number (12 objects) of detections. In addition, no correlation between the disc-averaged D/H and 14N/15N ratios have been measured by Guzmán et al. (2017) in six protoplanetary discs, but again the conclusion is based on poor statistics. Another intriguing result is that the D- and 15N-enhancements are not always observed in the same place in pristine solar system material (Busemann et al. 2006; Robert & Derenne 2006). It is therefore important to gather more data in sources that are good candidates to represent the environment in which our Sun was born to place stringent constraints on current chemical models. In this respect, intermediate- and high-mass star-forming cores are interesting targets because growing evidence shows that the Sun was born in a rich cluster that also contained massive stars (Adams 2010; Banerjee et al. 2016). Moreover, Taquet et al. (2016) have recently proposed that the proto-Sun was born in an environment that was denser and warmer than has commonly been considered for a solar system progenitor. In any case, because the statistics is still poor, more observations in star-forming cloud cores in different evolutionary stages are useful for a better understanding of whether and how the 15N fractionation process might be influenced by evolution. An example of existing work about 15N fractionation in high-mass dense cores is Adande & Ziurys (2012). These authors worked with a larger beam than was used in our observations, and they did not have an evolutionary classification of the sources, which makes interpreting their results in an evolutionary study difficult.

We here report the first measurements of the 14N/15N ratio derived from HCN and HNC in a sample of 27 dense cores associated with different stages of the high-mass star formation process that have previously been studied for their deuterated molecules and the 15N-bearing species of N2H+ (Fontani et al. 2011, 2015a,b). In particular, Fontani et al. (2015a) for the first time measured the 14N/15N isotopic ratio in N2H+ towards the same sources. These new data therefore allow us to investigate the possible difference between nitrogen hydrides and nitrile-bearing species that have been proposed by both theoretical studies (Wirström et al. 2012) and observational findings (Hily-Blant et al. 2013a). We also report measurements of the D/H ratios for HNC to search for relations between the two isotopic ratios.

2. Observations

We performed observations of the J = 1–0 rotational transition of H15NC, HN13C, HC15N, and H13CN towards the 27 sources observed by Fontani et al. (2015a) from 6 to 9 June 2015, using the 3 mm receiver of the IRAM-30 m telescope. We refer to this paper for the description of the source sample. We simultaneously observed the J = 2–1 transition of DNC with the 2 mm receiver. Table 1 presents the observed spectral windows and some main technical observational parameters. Table 2 presents the hyperfine frequencies of H13CN. The atmospheric conditions were very stable during the whole observing period, with precipitable water vapour usually in the range 3–8 mm. The observations were made in wobbler-switching mode with a wobbler throw of 240′′. Pointing was checked almost every hour on nearby quasars, planets, or bright Hii regions. The data were calibrated with the chopper wheel technique (see Kutner & Ulich 1981), with a calibration uncertainty of about 10% . The spectra were obtained in antenna temperature units, TA∗\hbox{$T_{\rm A}^{*}$}, and then converted into main beam brightness temperature, TMB, via the relation TA∗=TMB(Beff/Feff)=TMBηMB\hbox{$T_{\rm A}^{*}=T_{\rm MB}(B_{\rm eff}/F_{\rm eff})=T_{\rm MB} \eta_{\rm MB}$}, where ηMB = Beff/Feff is the ratio between the main beam efficiency (Beff) of the telescope and the forward efficiency of the telescope (Feff). The spectra were obtained with the fast Fourier transform spectrometers with the finest spectral resolution (FTS50), providing a channel width of 50 kHz. All calibrated spectra were analysed using the Gildas1 software developed at the IRAM and the Observatoire de Grenoble. Baselines in the spectra were all fitted by constant functions or polynomials of order 1. The rest frequencies used for the line identification were taken from different laboratory works: HC15N from Cazzoli et al. (2005), H13CN from Cazzoli & Puzzarini (2006), HN13C from van der Tak et al. (2009), H15NC from Pearson et al. (1976), and DNC from Bechtel et al. (2006). The other spectroscopic parameters used to derive the column densities were taken from the Cologne Molecular Database for Spectroscopy2 (CDMS; Müller et al. 2001, 2005), except for H15NC, for which we have used the Jet Propulsion Laboratory database3.

Table 1

Line rest frequencies and observational parameters.

Table 2

Frequencies of the hyperfine components of the transition H13CN(1–0).

3. Results

3.1. H15NC, HN13C, HC15N, and H13CN

The H15NC(1–0) line was detected in 26 cores (96.3%): 11 HMSCs, 8 HMPOs and 7 UC HIIs; the HC15N(1–0) line was detected in 24 cores (88.8%): 8 HMSCs, 9 HMPOs and 7 UC HIIs; the HN13C(1–0) line was detected in all sources (100%), and the H13CN(1–0) line was also clearly detected in all sources (100%).

To evaluate the isotopic ratios, we first of all used the 13C-bearing species of HCN and HNC because the main isotopologues are usually optically thick (e.g. Padovani et al. 2011). Then, we derived the 14N/15N ratio by correcting for the 12C/13C ratio. This latter was derived from the relation between this ratio and the source Galactocentric distance found for CN by Milam et al. (2005). Galactocentric distances were taken from Fontani et al. (2014, 2015a).

Neither H15NC(1–0) nor HC15N(1–0) have hyperfine structure, therefore all the lines were fitted with a single Gaussian. Conversely, HN13C(1–0) and H13CN(1–0) do have hyperfine structure. This cannot be resolved for the HN13C(1–0) because the line widths are always comparable to (or larger than) the separation in velocity of the hyperfine components. Then all the lines were also fitted with single Gaussians, and since the column densities are derived from the total integrated area of the rotational line, this simplified approach does not affect our measurements as long as the lines are optically thin, which is discussed in this section. On the other hand, it can overestimate the intrinsic line width. To estimate by how much, we fitted a line both with the Gaussian and the hyperfine method, and we found that with a single Gaussian, the line width is about 30% larger than is obtained with the hyperfine structure.

Table 3

Total column densities (beam averaged), computed as explained in Sect. 3, of H15NC, HN13C, and DNC.

Table 4

Total column densities (beam averaged), computed as explained in Sect. 3, of HC15N and H13CN(1–0) transitions.

Moreover, the fit using the hyperfine method demonstrated that the lines are optically thin. Finally, for H13CN(1–0), for which we were able to resolve the hyperfine structure, we fitted the components simultaneously, assuming that they have the same excitation temperature (Tex) and line width, and that the separation in velocity is fixed to the laboratory value. This fitting procedure gave good results in all spectra. The Gaussian fitting results are listed in Tables A.1 and A.2, and the hyperfine fitting results are listed in Table A.3. In Figs. B.1–B.4 we show the spectra of the H15NC, HC15N, HN13C, and H13CN(1-0) lines for all the 27 sources. The hyperfine components of H13CN(1–0) were always detected. The spectra of AFGL5142-EC may be partially contaminated from the nearby core AFGL5142-mm, but this emission is not expected to be dominant because the angular separation between the two cores is 30′′ (Busquet et al. 2011), the same as the beam of the telescope.

The total column densities of the four species averaged within the beam were evaluated from the total line-integrated intensity using Eq. (A4) of Caselli et al. (2002), which assumes that Tex is the same for all transitions within the same molecule, and for optically thin conditions. The assumption of optically thin lines is justified because from all the hyperfine fits of H13CN(1–0) we find τ ≪1 and well-constrained (Δτ/τ ≤ 1/3). We assume that the lines of the other isotopologues are also optically thin. We assume local thermal equilibrium (LTE) conditions, as all the observed sources have average H2 volume densities of the order of ~ 105 cm-3 (see Fontani et al. 2011, and references therein), that is, comparable to or marginally lower than the critical densities of the observed lines, thus this assumption is also reasonable. Because Tex cannot be deduced from our optically thin spectra (and also because we have only one transition), we adopted as Tex the kinetic temperatures given by Fontani et al. (2011), who derived them following the method described in Tafalla et al. (2004) based on Monte Carlo models, from which they obtained a relation among the kinetic temperature, Tk, and the NH3 rotation temperature between metastable levels. The Tk values are given in the last column of Tables 3 and 4. This last assumption is critical for the single column densities, but the HN13C/H15NC and H13CN/HC15N column density ratios do not change significantly by varying Tk between 20 and 100 K: changes are of one per cent or of ten percent, depending on the source. All the column densities and the parameters used to derive them (line integrated intensities) are given in Tables 3 and 4. We consider lines with TMBpeak≥3σ\hbox{$T_{\rm MB}^{\rm peak}\geq 3\sigma$} as detections. For the lines that are not clearly detected, we distinguished between those with 2.5σ≤TMBpeak<3σ\hbox{$2.5\sigma \leq T_{\rm MB}^{\rm peak}< 3\sigma$}, and those with TMBpeak<2.5σ\hbox{$T_{\rm MB}^{\rm peak}<2.5\sigma$}. For the first lines, considered as tentative detections, we computed the total column densities as explained above, and for the latter lines we give an upper limit to the integrated areas and hence to the total column densities using

∫TMB dv=Δv1/2TMBpeak2ln2π,$$ \int T_{\rm MB}\,{\rm d}v=\frac{\Delta v_{1/2}T_{\rm MB}^{\rm peak}}{2\sqrt{\frac{\rm ln2}{\pi}}}, $$where σ is the rms of the spectra, TMBpeak\hbox{$T_{\rm MB}^{\rm peak}$} is taken equal to 3σ, and Δv1/2 is the average value of the FWHM of the lines that are clearly detected for the corresponding transition and evolutionary stage of the source. The average value Δv1/2 for the high-mass starless cores in our data is ΔvHMSC = 2.2 ± 0.3 km s-1, while for the high-mass protostellar objects, it is ΔvHMPO = 1.8 ± 0.2 km s-1. Finally, we have derived the column density uncertainties from the errors on the line areas for optically thin lines, given by σ×Δv×N\hbox{$\sigma\times\Delta v\times\sqrt{N}$} (σ = rms noise in the spectrum, Δv = spectral velocity resolution, and N = number of channels with signal) and taking into account the calibration error (10%) for the TMB. Conversely, uncertainties in the 14N/15N ratios were computed from the propagation of errors on the column densities, as explained above, without taking the calibration uncertainties into account because the lines were observed in the same spectrum (see Sect. 2), so that the calibration error cancels out in their ratio.

thumbnail Fig. 1

Top panels: column density of HCN compared with that of HC15N (left) and of HNC compared with that of H15NC (right). Bottom panels: comparison between the 14N/15N isotopic ratios derived from the column density ratios N(HCN)/N(HC15N) and N(HNC)/N(H15NC) (left) and comparison between the H/D and 14N/15N isotopic ratios in HNC (right). In all panels, the filled circles represent the detected sources (black = HMSCs; red = HMPOs; blue = UC HIIs). The open triangles in the top panels are the upper limits on N(H15NC) (right) and N(HC15N) (left), while in the bottom panels the open triangles indicate lower limits on either N(HCN)/N(HC15N) or N(HNC)/N(H15NC). The filled squares represent tentative detections. The solid lines in the top panels indicate the mean atomic composition as measured in the terrestrial atmosphere (TA) and in the protosolar nebula (PSN), while in the bottom left panel, the solid line indicates the locus of points where N(HCN)/N(HC15N) is equal to N(HNC)/N(H15NC).

3.2. DNC

We also detected the rotational transition DNC(2–1) for all the 27 sources. These lines are used to measure the D/H ratio, which is compared with the 14N/15N ratio. Such a high detection rate indicates that deuterated gas is present at every stage of the massive star and star cluster formation process, as has been noted by Fontani et al. (2011). The transition possesses a hyperfine structure that is not resolved because of the broad line widths (see Table A.2). Therefore the lines where fitted with a single Gaussian. To estimates how much the line widths are overestimated using Gaussian fits, we fitted a line both with the Gaussian and the hyperfine method, and we found that with a single Gaussian, the line width is about 10% larger than that obtained with the hyperfine structure. We point out that for evolved sources (HMPOs and UC HIIs) another line in the spectra partly overlaps DNC(2-1). This is identified as acetaldehyde at 152.608 GHz (the JKa,Kc = 80,8−70,7 transition). The fact that the line is only detected in the evolved objects is consistent with the idea that acetaldehyde is probably released from grain mantles because it is detected only in warmer and more turbulent objects (see e.g. Codella et al. 2015). When we fitted the DNC lines with Gaussians, we excluded the contribution of this line by fitting the two lines simultaneously, and when possible, we excluded this line from the fit. The fitting results are listed in Table A.2.

We determined the column densities under the assumption of optically thin conditions and the same Tex for all transitions using Eq. (A4) of Caselli et al. (2002) in this case as well,NTOT=8πνij3c3Aij1giW(Jνij(Tex)−Jνij(TBG))1(1−exp(−hνijkTex))Q(Tex)exp(−EjkTex),\begin{equation} N_{\rm TOT}\!=\!\!\frac{8\pi\nu_{ij}^{3}}{c^{3}A_{ij}}\frac{1}{g_{i}}\frac{W}{(J_{\nu_{ij}}(T_{\rm ex})\!-\!\!J_{\nu_{ij}}(T_{\rm BG}))}\frac{1}{\left(1\!\!-\!\exp{(-\frac{h\nu_{ij}}{kT_{\rm ex}})}\right)}\frac{Q(T_{\rm ex})}{\exp{\left(-\frac{E_{j}}{kT_{\rm ex}}\right)}} , \end{equation}(1)where νij is the transition frequency, Aij is the Einstein coefficient of spontaneous emission, gi is the statistical weight of the upper level, Ej is the energy of the lower level, c is the speed of light, k is the Boltzmann constant, Tex is the excitation temperature of the transition, Q(Tex) is the partition function at temperature Tex, and TBG is the background temperature (2.7 K), Jνij(T) is

Jνij(T)=hνijk1exp(hνijkT)−1,$$ J_{\nu_{ij}}(T)=\frac{h\nu_{ij}}{k}\frac{1}{\exp{(\frac{h\nu_{ij}}{kT})}-1}, $$W is the integrated intensity of the line (TB∫ dv=(∫TMB dv)ΩMBΩs\hbox{$\int T_{\rm B}\,{\rm d}v=(\int T_{\rm MB}\,{\rm d}v)\frac{\Omega_{\rm MB}}{\Omega_{\rm s}}$}, where TB is the brightness temperature, ΩMB is the solid angle of the main beam, and Ωs is the solid angle of the source). However, in this case, the HNC/DNC ratio depends on the temperature because of the different excitation conditions of the two transitions observed ((2–1) for DNC and (1–0) for HN13C), so that the ratio depends on the temperature by the factor exp(Ej/kT). We also corrected the HNC/DNC ratios for the different beams of the antenna at the frequencies of the two lines, and we assumed that the emissions of DNC(2–1) and HNC(1–0) are less extended than the beam size of DNC(2–1). This last correction results in a factor 3.09 that is to be multiplied with the HNC/DNC ratio:

(1.22λ1D1.22λ2D)2=(λ1λ2)2=(ν2ν1)2=3.09,$$ \left(\frac{1.22\frac{\lambda_{1}}{D}}{1.22\frac{\lambda_{2}}{D}}\right)^{2}=\left(\frac{\lambda_{1}}{\lambda_{2}}\right)^{2}=\left(\frac{\nu_{2}}{\nu_{1}}\right)^{2}=3.09, $$where λ1 and λ2 are the wavelengths of the HNC(1–0) and DNC(2–1) transitions, respectively (and ν1 and ν2 are the corresponding frequencies). In Fig. B.5 we show the spectra of DNC(2–1) for all the 27 sources. The total column densities are listed in Table 3. Finally, we derived the errors as explained in Sect. 3.1, but here for the D/H ratios we also considered the calibration uncertainties because the two lines were observed in separate setups.

4. Discussion

4.1. 15N-fractionation

The comparison between the column densities of the 15N-containing species and those of their main isotopologues, derived as explained in Sect. 3, is shown in Fig. 1. The corresponding 14N/15N ratios are given in Tables 3 and 4.

We first discuss the 15N-fractionation found for HCN. The top left panel in Fig. 1 shows no large spread of measured values, which are very similar to the value found for the PSN. The mean values for the three evolutionary stages are 346±37 for HMSCs, 363±25 for HMPOs, and 369±25 for UC HIIs. Although the HMSCs have the highest 15N-enrichment, the mean values for the three evolutionary categories are however consistent within the errors, which indicates that time does not seem to play a role in the fractionation of nitrogen (at least until the formation of a Hii region). In the top panels of Fig. 2 we show the 14N/15N ratios calculated for HCN in the 27 sources as a function of the Galactocentric distance, of the line width, and of the kinetic temperature, respectively: again, there is no evidence of a trend of these ratios with any of the adopted parameters. In particular, the lack of correlation with either temperature or line width, which are both thought to increase with the evolution of the source, confirms the independence of the 14N/15N ratio from the core age. The lack of correlation with evolutionary parameters was also found by Fontani et al. (2015a) in N2H+, but the dispersion of the ratios is clearly much smaller in our study (from 180 up to 1300 in Fontani et al. 2015a, and from 250 up to 650 in this work). The reason may also be that they have larger uncertainties on average.

thumbnail Fig. 2

Top panels: N(HCN)/N(HC15N) as a function of Galactocentric distances, line widths, and kinetic temperatures. Bottom panels: N(HNC)/N(H15NC) as a function of Galactocentric distances, line widths, and kinetic temperatures. The symbols are the same as in the bottom panels of Fig. 1.

We now examine the 15N- fractionation found for HNC: we do not find a large spread of values here either compared with the results reported in Fontani et al. (2015a). The 14N/15N ratios are distributed within ~250 and 630 (top right panel in Fig. 1). The mean values for the three evolutionary stages are 428±40 for HMSCs, 462±31 for HMPOs, and 428±29 for UC HIIs, which is consistent within the errors with the 14N/15N ratio measured for the PSN of about 441 (from the solar wind). For this molecule, no trend between the isotopic ratio and either the line width or the kinetic temperature was found either (bottom panels of Fig. 2). In conclusion, time does not seem to play a role for either ratio.

The lack of correlation between the 14N/15N isotopic ratios and evolutionary parameters or physical parameters (FWHM, Tk) are somewhat consistent with the chemical model of Roueff et al. (2015), who predicted no fractionation of HCN and HNC in cold and dense conditions. However, their models are more appropriate for low-mass dense cores, with a Tk of 10 K, than for the warmer high-mass pre-stellar objects of the type studied in this work. This may explain the lack of carbon fractionation in our work, which is in contrast to their predictions. We stress that 13C may in theory have reduced abundances because nitriles and isonitriles are predicted to be significantly depleted in 13C (Roueff et al. 2015). However, this depletion is at most of a factor 2 (see Fig. 4 in Roueff et al. 2015) and is derived from a chemical model with a fixed kinetic temperature of 10 K, which is not the average kinetic temperature of our sources (not even of the starless cores, see Table 3). The predictions of this model may therefore not be appropriate for our objects. Furthermore, the 13C-fractionation is dependent on the time and temperature evolution (Szűcs et al. 2014; Röllig & Ossenkopf 2013), and observational tests to verify whether this theoretical effect is real have yet to be performed.

thumbnail Fig. 3

Top panel: column density of HNC compared with that of DNC. The three lines represent the mean values of D/H (see text) for different evolutionary stages: black for HMSCs, red for HMPOs and blue for UC HIIs. Bottom panels: N(HNC)/N(DNC) as a function of line widths and kinetic temperatures. The filled circles have the same meaning as in Fig. 1.

Intriguingly, when we performed a Kolmogorov-Smirnov test, which is a non-parametric statistical test, to determine whether the two data sets (the first set is the HC14N/HC15N ratios and the second is the H14NC/H15NC ratios) belong to the same distributions, we found a P-value = 0.078. The test indicates independence only if P ≤ 0.05, and this could mean that we cannot conclude anything about the dependence or the independence of the two samples with our results. It is known from the literature that H15NC and HC15N can form through the same reactions (see Terzieva & Herbst 2000; Wirström et al. 2012). We also checked the H15NC/HC15N column density ratio, and we found that it is roughly more than 1 for HMSCs (with an average value of ~ 1 ± 0.5) and roughly less than 1 for HMPOs and UC HIIs (with an average value of ~ 0.3 ± 0.1). Loison et al. (2014) showed that in cold regions, where C and CO are depleted, the dissociative recombination of HCNH+ acts to isomerize HCN into HNC, producing HCN/HNC ratios close to or slightly above one. In hot gas, the CN + H2 reaction with its high activation energy (2370 K; Jacobs et al. 1989) and the isomeration process C + HNC → C + HCN (Loison et al. 2014) may be responsible for the high HCN/HNC measured in hot cores (Schilke et al. 1992) and in young stellar objects (Schöier et al. 2002). That our HC15N/H15NC ratios are always close to one appears to imply that both species are probing similar low-temperature gas.

Interestingly, the 14N/15N ratio measured towards the source HMSC-G034F2 is an outlier (too high) in the distribution of the ratios in both HCN and HNC (but with larger errors than in other objects): a similarly high value was found by Bizzocchi et al. (2013) in L1544, which is a typical low-mass pre-stellar core. In particular, Bizzocchi et al. found a value of 14N/15N ≃ 1000± 200 from the N2H+. This result, and the results measured in HMSC-G034F2, are not consistent with the predictions of current models on nitrogen chemistry for N2H+, HCN, and HNC (Roueff et al. 2015; Hily-Blant et al. 2013a). The 14N/15N ratios in HCN and HNC for cores G028-C1, G034-C3, G028-F2, G034-F1, and G034-G2 have also been studied independently by Zeng et al. (2017), and both works obtain consistent results (within a factor of two). Taking into account HC15N in G028-C1 and comparing our spectrum with that of Zeng et al. (2017) at the same velocity resolution of 0.68 km s-1, we obtained an rms of 0.01 K, while they obtained an rms of 0.02 K. Our TMBpeak\hbox{$T_{\rm MB}^{\rm peak}$} is 0.03 K, that is, more than 2.5σ, so we can refer to this line as a tentative detection, while Zeng et al. (2017) reported an rms that is comparable with the peak of our line. They also derived the upper limit of the column density using a Δv1/2 of the line of 2 km s-1, while with our fit we obtained a Δv1/2 of 4 km s-1: this is the reason for the factor 2 of difference in the integrated area of this line in the two works. This effect is the same in the other spectra (compare Figs. A1–A3 of their work with our Figs. B1–B4), and the different integrated areas are also due to the different resolution and sensitivity.

4.2. D-fractionation

The top panel of Fig. 3 shows the comparison between the column densities of DNC and HNC, derived as explained in Sect. 3. The N(HNC)/N(DNC) mean values obtained for the three evolutionary stages are 823±94 for HMSCs, 983±189 for HMPOs, and 1275±210 for UC HIIs, indicated by three different lines in Fig. 3. Accordingly, the D/H mean ratios are (1.4±0.2) × 10-3 for HMSCs, (1.3±0.2) × 10-3 for HMPOs, and (0.9±0.1) × 10-3 for UC HIIs. The D/H mean values were obtained from computing the average for all the D/H values (Table 3). We note that the D/H values are slightly higher in the early stages, but because of the large dispersions, the differences between the three evolutionary categories are not statistically significant. This result confirms the marginally decreasing trend found by Fontani et al. (2014), derived from the DNC (1–0) transition in a subsample of the sources observed in this work. Fontani et al. (2014) found an average D/H of 0.012, 0.009, and 0.008 in HMSCs, HMPOs, and UC HIIs, respectively, with no statistically significant differences among the three evolutionary groups.

4.3. Comparison between D/H and 14N/15N

Considering the D- and N-fractionation for HNC towards the same sources, we can find an indication as to whether the two fractionations are linked. The bottom right panel in Fig. 1 shows HNC/DNC as a function of HNC/H15NC. The data show an independence between the two data sets, and this can be shown by computing the Kendall τ test. This is a non-parametric test used to measure the ordinal association between two data sets; its definition is

τ=(#concordantpairs)−(#discordantpairs)n(n−1)2,$$ \tau=\frac{(\text{$\#$ concordant pairs})-(\text{$\#$ discordant pairs})}{\frac{n(n-1)}{2}}, $$with n the number of the total pairs. If τ = 1, there is a full correlation, if τ = −1, there is full anti-correlation, and if τ = 0, the two data sets are independent. We chose this statistical test because compared to other non-parametric tests (i.e. the Spearman ρ correlation coefficient), this test is more robust; moreover, it allows us to compare our analysis with that performed by Fontani et al. (2015a), who used the same test. Following this, we computed τ for the two data sets HNC/DNC and HNC/H15NC, and we found τ ~ 0.13, which means that the D- and 15N- fractionation in these sources for HNC are independent. This result must be compared with that of Fontani et al. (2015a), which suggested a possible anti-correlation between the two isotopic ratios for N2H+. This finding arises mostly from the fact that the 14N/15N ratio does not vary with the core evolution, while the D-fractionation shows a faint decreasing trend. This result indicates that the parameters that cause D-enrichment in HCN and HNC may not influence the fractionation of nitrogen. The independence we found reflects what has been found in some pristine solar system material, in which the spots of high D-enrichments are not always spatially coincident with those with high 15N enrichment. From the point of view of the models, as discussed in Wirström et al. (2012), the D- and 15N-enrichments do not need to be spatially correlated (although they could be produced by the same mechanism, i.e. exothermic reactions due to the different zero-point energy of the heavier isotope), because relevant reactions for D- and 15N-enrichments have different energy barriers. Roueff et al. (2015) recently reviewed some of the reactions in 15N-fractionation, and based on the model, concluded that other modeling work is necessary to fully understand the relation of the two fractionation processes.

4.4. Comparison between N(HNC) and N(HCN)

For completeness, we also calculated the ratio between N(HNC) and N(HCN). Figure 4 shows the column density of HCN versus HNC, and we immediately note that HMSCs show values of HCN/HNC ≲ 1 unlike HMPOs and UC HIIs, for which the ratio is >1. The mean values are 0.8 ±0.2 for HMSCs, 2.5 ±0.2 for HMPOs, and 2.0±0.2 for UC HIIs.

thumbnail Fig. 4

Column density of HCN compared with that of HNC. The solid line is the locus where N(HCN)/N(HNC) is equal to 1, as statistically expected. The filled circles have the same meaning as in Fig. 1.

We computed the HNC/HCN ratio in the two data sets of HMSCs and HMPOs+UC HIIs. Using the non-parametric statistical Kolmogorov-Smirnov test, we obtain a P-value = 0.025, which indicates that the two distributions are different. Schilke et al. (1992) reported predictions and observations for the Orion hot core OMC-1, which is the prototype of a high-mass hot core. They found that the HCN/HNC abudance ratio is very high (~80) in the vicinity of Orion-KL, but it rapidly declines in positions adjacent to values of ~5. They compared the observations with the predictions of molecular cloud chemistry and found an agreement with steady-state models. More recently, Jin et al. (2015) found that the abundance ratio increases if the sources evolved from IRDCs to UC HIIs, and they suggested that this might occur for neutral-neutral reactions where HNC is selectively consumed for T ≳ 24 K (Hirota et al. 1998):

HNC + H → HCN + H ; $$ {\rm HNC+H} \rightarrow {\rm HCN + H}; $$

HNC+O→NH+CO.$$ {\rm HNC+O} \rightarrow {\rm NH +CO}. $$For the first reaction, a 2000 K barrier was found, and the second barrier has not yet been theoretically and experimentally studied, but if it were possible, N(HNC)/N(HCN) would decrease rapidly in warm gas. We therefore emphasize the importance of future laboratory measurements of these later reactions, especially at low temperatures.

thumbnail Fig. 5

14N/15N ratios as a function of Galactocentric distances (for HNC and HCN together). The symbols are the same as in the bottom panels of Fig. 1. The black line is the linear fit computed for the plotted data, and the grey line represents the fit of Adande & Ziurys (2012).

4.5. Comparison with the 14N/15N Galactocentric gradient

Adande & Ziurys (2012) evaluated the 14N/15N ratio across the Galaxy through millimetre observations of CN and HNC in regions of star formation. They also enlarged the sample with the high-mass sources observed by Dahmen et al. (1995) via HCN observations. In order to compare our results with their work, we have plotted our 14N/15N ratios (for HCN and HNC together) as a function of the Galactocentric distance of the sources (see Table C.1 and Fig. 5). We computed a linear fit and obtained the relation 14N/15N=(4.9±7.6)kpc-1×DGC+(366.5±65.6),\begin{equation} ^{14}\text{N}/^{15}\text{N}=(4.9\pm7.6) \text{ kpc}^{-1} \times D_{\rm GC} + (366.5\pm65.6) , \end{equation}(2)which disagrees with the relation obtained by Adande & Ziurys (2012), 14N/15N=(21.1±5.2)kpc-1×DGC+(123.8±37.1).\begin{equation} ^{14}\text{N}/^{15}\text{N}=(21.1\pm5.2) \text{ kpc}^{-1} \times D_{\rm GC} + (123.8\pm37.1). \end{equation}(3)This indicates that our data set does not confirm the gradient found by Adande & Ziurys (2012). We speculate that as the beam used by Adande & Ziurys (2012) is larger and hence more influenced by the diffuse surrounding material, the difference in ratios may be due to the material surrounding the target cores, in which this ratio may be lower. We also note that Adande & Ziurys (2012) derived their fit using the source SgrB2(NW), which is located at a Galactocentric radius of 0.1 kpc: for this source, they obtained a 14N /15N lower limit of 164. This value strongly influences their fit results and hence could have a strong effect on both the slope and the intercept with the y-axis. Any comparison between the two fit results must therefore be taken with caution.

5. Conclusions

We have observed the J = 1–0 rotational transitions of H15NC, HN13C, H13CN, and HC15N together with the J = 2–1 transition of DNC towards 27 massive star-forming cores in different evolutionary stages in order to derive the 15N- and D-fractionation, and to compare the ratios with each other. We find 14N/15N in HCN between ~200 and ~700, and in HNC between ~260 and ~600, with a small spread around the PSN value of 441. Comparing the 14N/15N ratios for different evolutionary stages, we did not find any trend, indicating that time does not seem to play a role in the N-fractionation; furthermore, we cannot conclude about the correlation between the 15N-fractionation for the two molecules HNC and HCN. Considering both D- and 15N-fractionation for HNC towards the same sources, we found no correlation. This is consistent with the lack of correlation found by Fontani et al. (2015a) in N2H+: the causes of D-enrichment in HCN and HNC do not affect the 15N-fractionation. Our findings are in agreement with the recent chemical models of Roueff et al. (2015). The independence between D/H and 14N/15N ratios confirms the recent findings of Guzmán et al. (2017) in protoplanetary discs. Conversely, the low values of 14N/15N found by them (and the tentative decrease with decreasing distance to the star) does not contradict our findings because our results are obtained on angular scales much larger than that of a protoplanetary disc. All this indicates that the 15N enrichment is a local effect that does not involve the larger-scale envelope.

1

The Gildas software is available at http://www.iram.fr/IRAMFR/GILDAS

Acknowledgments

We are grateful to the IRAM-30 m telescope staff for their help during the observations. Many thanks to the anonymous referee for the careful reading of the paper and the comments that improved the work. Paola Caselli acknowledges support from the European Research Council (project PALs 320620).

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Appendix A: Fit results

In this appendix we show the results of the fitting procedure for the HN13C(1–0), H15NC(1–0), HC15N(1–0), DNC(2–1), and H13CN(1–0) lines of all sources (see Sects. 3.1 and 3.2).

Table A.1

Values obtained with a Gaussian fit to the HN13C and H15NC(1–0) lines.

Table A.2

Values obtained with Gaussian fits to the HC15N and DNC(2–1) lines.

Table A.3

Values obtained from the hyperfine fit of the H13CN(1–0) lines.

Appendix B: Spectra

In this appendix we show all spectra of the HN13C(1–0), H13CN(1–0), H15NC(1–0), HC15NC(1–0), and DNC(2–1) transitions for all the sources.

thumbnail Fig. B.1

Spectra of HN13C(1–0) obtained for the sources classified as HMSCs (first column), HMPOs (second column), and UC HII regions (third column). For each spectrum the x-axis represents a velocity interval of ± 20 km s-1 around the systemic velocity listed in Table 1 of Fontani et al. (2015b). The y-axis shows the intensity in main-beam temperature units. The red curves are the best Gaussian fits obtained with CLASS. For some sources (I20293-WC, G034-G2, G034-F2, G5.89-0.39, G034-F1, G028-C1, and I23385), we have observed two components: we have fitted both lines and used only the line centered on the systemic velocity of the source to compute the column densities.

thumbnail Fig. B.2

Same as Fig. A.3 for H13CN(1–0). Here the red curves are the best hyperfine fits obtained with CLASS. We note the presence of the second velocity component in the same sources that were indicated in the caption of Fig. B.1.

thumbnail Fig. B.3

Same as Fig. A.3 for H15NC(1–0). For I21307, this line was not detected, and we have obtained a column density upper limit, as explained in Sect. 3.1.

thumbnail Fig. B.4

Same as Fig. A.3 for HC15N(1–0). For G034-G2, G028-C3, and G034-F2, the line was not detected, and we obtained a column density upper limit, as explained in Sect. 3.1. For G034-F1 we did not have a clear detection, but we have obtained from the Gaussian fit that TMBpeak≳2.5σ\hbox{$T_{\rm MB}^{\rm peak}\gtrsim2.5\sigma$}, we have computed as usual the column density, and we refer to this latter as “tentative detection” (see also Table A.2).

thumbnail Fig. B.5

Same as Fig. A.3 for DNC(2–1). For evolved sources (HMPOs and UC HII) the acetaldehyde line is also present, as explained in Sect. 3.2: when possible, we have excluded the line from the Gaussian fit, otherwise we have fitted the two lines together and used only the DNC(2–1) to compute the column densities.

Appendix C: Comparison with other molecules

In this appendix, we compare the 14N/15N ratios obtained in this paper and in Fontani et al. (2015a) for all the common sources. The 12C/13C we used and the Galactocentric distances are also listed.

Table C.1

Comparison between 14N/15N ratios obtained from HNC and HCN in this work and those obtained from N2H+ and CN by Fontani et al. (2015a).

All Tables

Table 1

Line rest frequencies and observational parameters.

In the text
Table 2

Frequencies of the hyperfine components of the transition H13CN(1–0).

In the text
Table 3

Total column densities (beam averaged), computed as explained in Sect. 3, of H15NC, HN13C, and DNC.

In the text
Table 4

Total column densities (beam averaged), computed as explained in Sect. 3, of HC15N and H13CN(1–0) transitions.

In the text
Table A.1

Values obtained with a Gaussian fit to the HN13C and H15NC(1–0) lines.

In the text
Table A.2

Values obtained with Gaussian fits to the HC15N and DNC(2–1) lines.

In the text
Table A.3

Values obtained from the hyperfine fit of the H13CN(1–0) lines.

In the text
Table C.1

Comparison between 14N/15N ratios obtained from HNC and HCN in this work and those obtained from N2H+ and CN by Fontani et al. (2015a).

In the text

All Figures

thumbnail Fig. 1

Top panels: column density of HCN compared with that of HC15N (left) and of HNC compared with that of H15NC (right). Bottom panels: comparison between the 14N/15N isotopic ratios derived from the column density ratios N(HCN)/N(HC15N) and N(HNC)/N(H15NC) (left) and comparison between the H/D and 14N/15N isotopic ratios in HNC (right). In all panels, the filled circles represent the detected sources (black = HMSCs; red = HMPOs; blue = UC HIIs). The open triangles in the top panels are the upper limits on N(H15NC) (right) and N(HC15N) (left), while in the bottom panels the open triangles indicate lower limits on either N(HCN)/N(HC15N) or N(HNC)/N(H15NC). The filled squares represent tentative detections. The solid lines in the top panels indicate the mean atomic composition as measured in the terrestrial atmosphere (TA) and in the protosolar nebula (PSN), while in the bottom left panel, the solid line indicates the locus of points where N(HCN)/N(HC15N) is equal to N(HNC)/N(H15NC).
In the text
thumbnail Fig. 2

Top panels: N(HCN)/N(HC15N) as a function of Galactocentric distances, line widths, and kinetic temperatures. Bottom panels: N(HNC)/N(H15NC) as a function of Galactocentric distances, line widths, and kinetic temperatures. The symbols are the same as in the bottom panels of Fig. 1.
In the text
thumbnail Fig. 3

Top panel: column density of HNC compared with that of DNC. The three lines represent the mean values of D/H (see text) for different evolutionary stages: black for HMSCs, red for HMPOs and blue for UC HIIs. Bottom panels: N(HNC)/N(DNC) as a function of line widths and kinetic temperatures. The filled circles have the same meaning as in Fig. 1.
In the text
thumbnail Fig. 4

Column density of HCN compared with that of HNC. The solid line is the locus where N(HCN)/N(HNC) is equal to 1, as statistically expected. The filled circles have the same meaning as in Fig. 1.
In the text
thumbnail Fig. 5

14N/15N ratios as a function of Galactocentric distances (for HNC and HCN together). The symbols are the same as in the bottom panels of Fig. 1. The black line is the linear fit computed for the plotted data, and the grey line represents the fit of Adande & Ziurys (2012).

In the text
thumbnail Fig. B.1

Spectra of HN13C(1–0) obtained for the sources classified as HMSCs (first column), HMPOs (second column), and UC HII regions (third column). For each spectrum the x-axis represents a velocity interval of ± 20 km s-1 around the systemic velocity listed in Table 1 of Fontani et al. (2015b). The y-axis shows the intensity in main-beam temperature units. The red curves are the best Gaussian fits obtained with CLASS. For some sources (I20293-WC, G034-G2, G034-F2, G5.89-0.39, G034-F1, G028-C1, and I23385), we have observed two components: we have fitted both lines and used only the line centered on the systemic velocity of the source to compute the column densities.

In the text
thumbnail Fig. B.2

Same as Fig. A.3 for H13CN(1–0). Here the red curves are the best hyperfine fits obtained with CLASS. We note the presence of the second velocity component in the same sources that were indicated in the caption of Fig. B.1.
In the text
thumbnail Fig. B.3

Same as Fig. A.3 for H15NC(1–0). For I21307, this line was not detected, and we have obtained a column density upper limit, as explained in Sect. 3.1.
In the text
thumbnail Fig. B.4

Same as Fig. A.3 for HC15N(1–0). For G034-G2, G028-C3, and G034-F2, the line was not detected, and we obtained a column density upper limit, as explained in Sect. 3.1. For G034-F1 we did not have a clear detection, but we have obtained from the Gaussian fit that TMBpeak≳2.5σ\hbox{$T_{\rm MB}^{\rm peak}\gtrsim2.5\sigma$}, we have computed as usual the column density, and we refer to this latter as “tentative detection” (see also Table A.2).
In the text
thumbnail Fig. B.5

Same as Fig. A.3 for DNC(2–1). For evolved sources (HMPOs and UC HII) the acetaldehyde line is also present, as explained in Sect. 3.2: when possible, we have excluded the line from the Gaussian fit, otherwise we have fitted the two lines together and used only the DNC(2–1) to compute the column densities.
In the text

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