First interferometric study of enhanced N-fractionation in N$_{2}$H$^{+}$: the high-mass star-forming region IRAS 05358+3543

TL;DR

This study uses high-resolution radio observations to investigate nitrogen isotope ratios in a high-mass star-forming region, revealing local environmental effects and potential gaps in current chemical models.

Contribution

First interferometric measurements of N-fractionation in a high-mass star-forming region, showing spatial variation and suggesting missing chemical processes in models.

Findings

N-14/N-15 ratios range from ~100 to ~220.

Local environmental effects influence N-fractionation.

Potential chemical abundance differences in protocluster cores.

Abstract

Nitrogen (N) fractionation is used as a tool to search for a link between the chemical history of the Solar System and star-forming regions. A large variation of $^{14}$N/$^{15}$N is observed towards different astrophysical sources, and current chemical models cannot reproduce it. With the advent of high angular resolution radiotelescopes it is now possible to search for N-fractionation at core scales. We present IRAM NOEMA observations of the J=1-0 transition of N$_{2}$H$^{+}$, $^{15}$NNH$^{+}$ and N$^{15}$NNH$^{+}$ towards the high-mass protocluster IRAS 05358+3543. We find $^{14}$N/$^{15}$N ratios that span from $\sim$100 up to $\sim$220 and these values are lower or equal than those observed with single-dish observations towards the same source. Since N-fractionation changes across the studied region, this means that it is regulated by local environmental effects. We find also the possibility, for one of the four cores defined in the protocluster, to have a more abundant $^{15}$NNH$^{+}$ with respect to N$^{15}$NNH$^{+}$. This is another indication that current chemical models may be missing chemical reactions or may not take into account other mechanisms, like photodissociation or grain surface chemistry, that could be important.

Tables

Table 1. Table 1: Coordinates of the dust continuum sources as observed by Beuther et al. ( 2007 ).

Source	RA[J2000]	DEC[J2000]
	(h m s)	($°$ $\prime$ $″$)
mm1a	05:39:13.08	35:45:51.3
mm1b	05:39:13.13	35:45:50.8
mm2	05:39:12.76	35:45:51.3
mm3	05:39:12.50	35:45:54.9
mm4a	05:39:12.46	35:45:40.9

Table 2. Table 2: Molecular transitions of the J=1–0 transition of N 2 H + , 15 NNH + and N 15 NH + .

Molecule	Frequency	Transition	$E_U$	log10(A${}_{}{}^a{}_{\mathrm{ij}}{}^{}$)
	(GHz)		(K)
N2H+b	93.1734	J=1–0	4.5	-4.44034
15NNH+	90.2635	J=1–0, F=1–1	4.3	-4.48168
	90.2639	J=1–0, F=2–1
	90.2645	J=1–0, F=0–1
N15NH+	91.2043	J=1–0, F=1–1	4.4	-4.46810
	91.2060	J=1–0, F=2–1
	91.2086	J=1–0, F=0–1

Table 3. Table 3: Observational parameters.

Spw	Synthesized beam	PA	$\Delta$v	rms
	($″\times ″$)	(°)	(km s-1)	(mJy/beam)
Continuum	3.01 $\times$ 2.61	-110.18	–	0.2
N2H+	3.04 $\times$ 2.64	-109.5	6.45	0.5
15NNH+	3.12 $\times$ 3.08	-41.91	0.5	7
N15NH+	3.12 $\times$ 3.08	-41.91	0.5	7

Table 4. Table 4: 14 N/ 15 N ratios derived from 15 NNH + towards P1a, P2a, P3a and P4a (left panel) and from N 15 NH + towards P1b, P2b, P3b and P4b (right panel).

Source	$T_{\mathrm{ex}}$	$\frac{N_2{H}^{+}}{{}_{}{}^{15}\mathrm{NNH}_{}^{+}}$	Source	$T_{\mathrm{ex}}$	$\frac{N_2{H}^{+}}{N^{15}{\mathrm{NH}}^{+}}$
	(K)			(K)
P1a	20	200$\pm$34	P1b	20	200$\pm$25
	30	186$\pm$31		30	217$\pm$29
	40	178$\pm$29		40	212$\pm$27
	50	173$\pm$27		50	222$\pm$30
P2a	20	180$\pm$18	P2b	20	129$\pm$17
	30	171$\pm$17		30	120$\pm$15
	40	156$\pm$16		40	108$\pm$13
	50	154$\pm$15		50	113$\pm$14
P3a	20	120$\pm$18	P3b	20	100$\pm$15
	30	114$\pm$17		30	100$\pm$16
	40	111$\pm$16		40	109$\pm$17
	50	109$\pm$15		50	100$\pm$15
P4a	20	117$\pm$14	P4b	20	187$\pm$33
	30	112$\pm$14		30	182$\pm$31
	40	120$\pm$15		40	185$\pm$33
	50	108$\pm$12		50	185$\pm$32

Table 5. Table 5: 14 N/ 15 N ratios derived from 15 NNH + and N 15 NH + towards I1, I2, I3 and I4.

Source	$T_{\mathrm{ex}}$	$\frac{N_2{H}^{+}}{{}_{}{}^{15}\mathrm{NNH}_{}^{+}}$		$\frac{N_2{H}^{+}}{N^{15}{\mathrm{NH}}^{+}}$
	(K)
I1	20	220$\pm$32		275$\pm$50
	30	214$\pm$30		250$\pm$41
	40	200$\pm$27		257$\pm$46
	50	200$\pm$28		244$\pm$40
I2	20	167$\pm$17		143$\pm$18
	30	150$\pm$15		120$\pm$15
	40	136$\pm$13		125$\pm$16
	50	146$\pm$14		127$\pm$16
I3	20	100$\pm$13		86$\pm$12
	30	114$\pm$16		89$\pm$13
	40	111$\pm$15		91$\pm$14
	50	109$\pm$15		86$\pm$12
I4	20	100$\pm$14		140$\pm$24
	30	100$\pm$14		150$\pm$29
	40	100$\pm$15		137$\pm$25
	50	93$\pm$14		130$\pm$23

Table 6. Table 6: 14 N/ 15 N ratios derived from 15 NNH + and N 15 NH + for D1, D2 and D3 (upper panel). In the lower panel the 14 N/ 15 N ratios derived from both 15 NNH + and N 15 NH + towards mm1 and mm3 in regions equivalent to the IRAM-30m beam are shown.

Source	$T_{\mathrm{ex}}$	$\frac{N_2{H}^{+}}{{}_{}{}^{15}\mathrm{NNH}_{}^{+}}$	$\frac{N_2{H}^{+}}{N^{15}{\mathrm{NH}}^{+}}$
	(K)
D1	20	$\ge$245	$\ge$204
	30	$\ge$231	$\ge$188
	40	$\ge$242	$\ge$200
	50	$\ge$250	$\ge$204
D2	20	336$\pm$96	$\ge$154
	30	327$\pm$91	$\ge$148
	40	316$\pm$87	$\ge$143
	50	292$\pm$77	$\ge$140
D3	20	$\ge$243	$\ge$340
	30	$\ge$250	$\ge$321
	40	$\ge$261	$\ge$353
	50	$\ge$250	$\ge$333
Source	$T_{\mathrm{ex}}$	$\frac{N_2{H}^{+}}{{}_{}{}^{15}\mathrm{NNH}_{}^{+}}$	$\frac{N_2{H}^{+}}{N^{15}{\mathrm{NH}}^{+}}$
	(K)
mm1	43	219$\pm$22	180$\pm$33
mm3	18	217$\pm$23	192$\pm$32

Table 7. Table 7: Values obtained with the fitting procedure described in Sect.  3.3 to the J=1–0 transition of N 2 H + and 15 NNH + towards P1a, P2a, P3a and P4a (top panel of the table), and of the J=1–0 transition of N 2 H + and N 15 NH + towards P1b, P1b, P3b and P4b (bottom panel of the table).

Source	N2H+(1–0)					15NNH+(1–0)
	$T_{\mathrm{ex}}$	$FWH{M}^a$	$v_{\mathrm{LSR}}$	$\tau$	$N_{\mathrm{tot}}$	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1014cm-2)	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)
P1a	20	11.1$\pm$0.6	-15.2$\pm$0.3	0.26$\pm$0.02	1.0$\pm$0.1	20	2.2$\pm$0.7	-16.2$\pm$0.3	0.002$\pm$0.001	5$\pm$1
	30	11.3$\pm$0.6	-15.2$\pm$0.3	0.16$\pm$0.02	1.3$\pm$0.2	30	2.2$\pm$0.7	-16.2$\pm$0.3	0.002$\pm$0.001	7$\pm$2
	40	11.3$\pm$0.6	-15.2$\pm$0.3	0.112$\pm$0.008	1.6$\pm$0.2	40	2.3$\pm$0.7	-16.2$\pm$0.3	0.002$\pm$0.001	9$\pm$2
	50	11.4$\pm$0.6	-15.2$\pm$0.3	0.087$\pm$0.006	1.9$\pm$0.3	50	2.3$\pm$0.7	-16.2$\pm$0.3	0.0013$\pm$0.0005	11$\pm$3
P2a	20	10.6$\pm$0.7	-14.7$\pm$0.3	0.25$\pm$0.02	0.9$\pm$0.1	20	1.8$\pm$0.3	-15.6$\pm$0.1	0.0027$\pm$0.0005	5$\pm$1
	30	10.8$\pm$0.7	-14.7$\pm$0.3	0.15$\pm$0.02	1.2$\pm$0.2	30	1.8$\pm$0.3	-15.6$\pm$0.1	0.0017$\pm$0.0003	7$\pm$1
	40	10.9$\pm$0.7	-14.7$\pm$0.3	0.105$\pm$0.008	1.4$\pm$0.2	40	1.8$\pm$0.3	-15.6$\pm$0.1	0.0013$\pm$0.0003	9$\pm$2
	50	10.9$\pm$0.7	-14.7$\pm$0.3	0.082$\pm$0.006	1.7$\pm$0.3	50	1.8$\pm$0.3	-15.6$\pm$0.1	0.0010$\pm$0.0002	11$\pm$2
P3a	20	11$\pm$1	-15.0$\pm$0.5	0.17$\pm$0.02	0.6$\pm$0.1	20	2.3$\pm$0.6	-16.6$\pm$0.2	0.002$\pm$0.001	5$\pm$1
	30	11$\pm$1	-15.0$\pm$0.5	0.10$\pm$0.02	0.8$\pm$0.1	30	2.3$\pm$0.6	-16.6$\pm$0.2	0.0013$\pm$0.0003	7$\pm$1
	40	11$\pm$1	-15.0$\pm$0.5	0.073$\pm$0.009	1.0$\pm$0.2	40	2.3$\pm$0.6	-16.6$\pm$0.2	0.0010$\pm$0.0002	9$\pm$2
	50	11$\pm$1	-15.0$\pm$0.5	0.057$\pm$0.007	1.2$\pm$0.2	50	2.3$\pm$0.6	-16.6$\pm$0.2	0.0008$\pm$0.0002	11$\pm$2
P4a	20	10.2$\pm$0.6	-14.6$\pm$0.3	0.20$\pm$0.02	0.7$\pm$0.1	20	2.7$\pm$0.6	-15.6$\pm$0.2	0.0021$\pm$0.0005	6$\pm$1
	30	10.3$\pm$0.6	-14.6$\pm$0.3	0.12$\pm$0.01	0.9$\pm$0.1	30	2.7$\pm$0.6	-15.6$\pm$0.2	0.0013$\pm$0.0003	8$\pm$2
	40	10.3$\pm$0.6	-14.6$\pm$0.3	0.089$\pm$0.007	1.2$\pm$0.2	40	2.7$\pm$0.6	-15.6$\pm$0.2	0.0009$\pm$0.0002	10$\pm$2
	50	10.4$\pm$0.6	-14.6$\pm$0.3	0.069$\pm$0.005	1.4$\pm$0.2	50	2.7$\pm$0.6	-15.6$\pm$0.2	0.0007$\pm$0.0002	13$\pm$3
Source	N2H+(1–0)					N15NH+(1–0)
	$T_{\mathrm{ex}}$	$FWH{M}^a$	$v_{\mathrm{LSR}}$	$\tau$	$N_{\mathrm{tot}}$	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1014cm-2)	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)
P1b	20	11.2$\pm$0.7	-15.2$\pm$0.3	0.27$\pm$0.02	1.0$\pm$0.2	20	1.2$\pm$0.2	-16.27$\pm$0.07	0.004$\pm$0.001	5$\pm$1
	30	11.4$\pm$0.7	-15.2$\pm$0.3	0.16$\pm$0.02	1.3$\pm$0.2	30	1.2$\pm$0.2	-16.27$\pm$0.07	0.0022$\pm$0.0004	6$\pm$1
	40	11.5$\pm$0.7	-15.2$\pm$0.3	0.114$\pm$0.009	1.7$\pm$0.3	40	1.2$\pm$0.2	-16.27$\pm$0.07	0.0016$\pm$0.0003	8$\pm$2
	50	11.5$\pm$0.7	-15.2$\pm$0.3	0.089$\pm$0.007	2.0$\pm$0.3	50	1.2$\pm$0.2	-16.27$\pm$0.07	0.0013$\pm$0.0002	9$\pm$2
P2b	20	11.0$\pm$0.7	-14.7$\pm$0.3	0.23$\pm$0.02	0.9$\pm$0.1	20	1.9$\pm$0.3	-15.5$\pm$0.1	0.004$\pm$0.001	7$\pm$2
	30	11.1$\pm$0.6	-14.7$\pm$0.3	0.14$\pm$0.01	1.2$\pm$0.2	30	1.9$\pm$0.3	-15.5$\pm$0.1	0.0023$\pm$0.0004	10$\pm$2
	40	11.2$\pm$0.6	-14.7$\pm$0.3	0.110$\pm$0.007	1.4$\pm$0.2	40	1.9$\pm$0.3	-15.5$\pm$0.1	0.0016$\pm$0.0003	13$\pm$3
	50	11.2$\pm$0.6	-14.7$\pm$0.3	0.078$\pm$0.006	1.7$\pm$0.3	50	1.9$\pm$0.3	-15.5$\pm$0.1	0.0013$\pm$0.0002	15$\pm$3
P3b	20	11$\pm$1	-14.9$\pm$0.5	0.20$\pm$0.03	0.7$\pm$0.1	20	2.0$\pm$0.3	-16.5$\pm$0.2	0.003$\pm$0.001	7$\pm$1
	30	11$\pm$1	-14.9$\pm$0.5	0.12$\pm$0.02	0.9$\pm$0.2	30	2.0$\pm$0.3	-16.5$\pm$0.2	0.0019$\pm$0.0004	9$\pm$2
	40	11$\pm$1	-14.9$\pm$0.5	0.08$\pm$0.01	1.2$\pm$0.2	40	2.0$\pm$0.3	-16.5$\pm$0.2	0.0014$\pm$0.0003	11$\pm$3
	50	11$\pm$1	-14.9$\pm$0.5	0.065$\pm$0.008	1.4$\pm$0.3	50	2.0$\pm$0.3	-16.5$\pm$0.2	0.0011$\pm$0.0002	14$\pm$3
P4b	20	10.7$\pm$0.6	-14.7$\pm$0.3	0.082$\pm$0.006	0.30$\pm$0.04	20	1.5$\pm$0.3	-16.7$\pm$0.2	0.0010$\pm$0.0003	1.6$\pm$0.4
	30	10.7$\pm$0.6	-14.7$\pm$0.3	0.051$\pm$0.004	0.40$\pm$0.06	30	1.5$\pm$0.3	-16.7$\pm$0.2	0.0006$\pm$0.0002	2.2$\pm$0.6
	40	10.7$\pm$0.6	-14.7$\pm$0.3	0.037$\pm$0.003	0.50$\pm$0.07	40	1.5$\pm$0.3	-16.7$\pm$0.2	0.0005$\pm$0.0001	2.7$\pm$0.7
	50	10.7$\pm$0.6	-14.7$\pm$0.3	0.029$\pm$0.002	0.61$\pm$0.09	50	1.5$\pm$0.3	-16.7$\pm$0.2	0.00040$\pm$0.00009	3.3$\pm$0.9

Table 8. Table 8: Values obtained with the fitting procedure described in Sect.  3.3 to the N 2 H + (1–0), 15 NNH + (1–0) and N 15 NH + (1–0) transitions towards I1, I2, I3 and I4.

Source	N2H+(1–0)
	$T_{\mathrm{ex}}$	$FWH{M}^a$	$v_{\mathrm{LSR}}$	$\tau$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1014cm-2)
I1	20	11.3$\pm$0.7	-15.3$\pm$0.3	0.29$\pm$0.03	1.1$\pm$0.2
	30	11.5$\pm$0.7	-15.3$\pm$0.3	0.17$\pm$0.02	1.5$\pm$0.2
	40	11.6$\pm$0.7	-15.3$\pm$0.3	0.124$\pm$0.009	1.8$\pm$0.3
	50	11.7$\pm$0.7	-15.3$\pm$0.3	0.096$\pm$0.007	2.2$\pm$0.3
I2	20	11.0$\pm$0.6	-14.7$\pm$0.3	0.25$\pm$0.02	1.0$\pm$0.2
	30	11.1$\pm$0.7	-14.7$\pm$0.3	0.15$\pm$0.02	1.2$\pm$0.2
	40	11.2$\pm$0.7	-14.7$\pm$0.3	0.109$\pm$0.009	1.5$\pm$0.3
	50	11.3$\pm$0.7	-14.7$\pm$0.3	0.084$\pm$0.007	1.9$\pm$0.3
I3	20	11.0$\pm$0.9	-15.0$\pm$0.5	0.17$\pm$0.02	0.6$\pm$0.1
	30	11.1$\pm$0.9	-15.0$\pm$0.5	0.10$\pm$0.01	0.8$\pm$0.2
	40	11.2$\pm$0.9	-15.0$\pm$0.5	0.073$\pm$0.008	1.0$\pm$0.2
	50	11.2$\pm$0.9	-15.0$\pm$0.5	0.057$\pm$0.006	1.2$\pm$0.2
I4	20	10.8$\pm$0.6	-14.8$\pm$0.3	0.18$\pm$0.02	0.7$\pm$0.1
	30	10.9$\pm$0.6	-14.8$\pm$0.3	0.111$\pm$0.008	0.9$\pm$0.1
	40	11.0$\pm$0.6	-14.7$\pm$0.3	0.079$\pm$0.006	1.1$\pm$0.2
	50	11.0$\pm$0.6	-14.7$\pm$0.3	0.062$\pm$0.004	1.3$\pm$0.2
Source	15NNH+(1–0)					N15NH+(1–0)
	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)
I1	20	1.7$\pm$0.5	-16.4$\pm$0.2	0.005$\pm$0.001	5$\pm$1	20	1.3$\pm$0.3	-16.1$\pm$0.1	0.005$\pm$0.001	4$\pm$1
	30	1.7$\pm$0.5	-16.4$\pm$0.2	0.002$\pm$0.001	7$\pm$2	30	1.3$\pm$0.3	-16.1$\pm$0.1	0.003$\pm$0.001	6$\pm$2
	40	1.7$\pm$0.5	-16.4$\pm$0.2	0.002$\pm$0.001	9$\pm$2	40	1.3$\pm$0.3	-16.1$\pm$0.1	0.002$\pm$0.001	7$\pm$2
	50	1.7$\pm$0.5	-16.4$\pm$0.2	0.002$\pm$0.001	11$\pm$2	50	1.3$\pm$0.3	-16.1$\pm$0.1	0.0019$\pm$0.0005	9$\pm$2
I2	20	1.6$\pm$0.4	-15.7$\pm$0.2	0.006$\pm$0.002	6$\pm$1	20	1.7$\pm$0.3	-15.5$\pm$0.1	0.007$\pm$0.001	7$\pm$2
	30	1.9$\pm$0.3	-15.6$\pm$0.1	0.0030$\pm$0.0005	8$\pm$2	30	1.7$\pm$0.3	-15.5$\pm$0.1	0.004$\pm$0.001	10$\pm$2
	40	1.9$\pm$0.3	-15.6$\pm$0.1	0.0022$\pm$0.0003	11$\pm$2	40	1.7$\pm$0.3	-15.5$\pm$0.1	0.003$\pm$0.001	12$\pm$3
	50	1.9$\pm$0.3	-15.6$\pm$0.1	0.0017$\pm$0.0003	13$\pm$2	50	1.7$\pm$0.3	-15.5$\pm$0.1	0.0024$\pm$0.0004	15$\pm$3
I3	20	2.6$\pm$0.6	-16.6$\pm$0.2	0.003$\pm$0.001	6$\pm$1	20	2.2$\pm$0.3	-16.5$\pm$0.2	0.005$\pm$0.001	7$\pm$2
	30	2.6$\pm$0.6	-16.6$\pm$0.2	0.002$\pm$0.001	7$\pm$2	30	2.2$\pm$0.3	-16.5$\pm$0.2	0.003$\pm$0.001	9$\pm$2
	40	2.6$\pm$0.6	-16.6$\pm$0.2	0.0015$\pm$0.0004	9$\pm$2	40	2.2$\pm$0.3	-16.5$\pm$0.2	0.0021$\pm$0.0004	11$\pm$3
	50	2.6$\pm$0.6	-16.6$\pm$0.2	0.0012$\pm$0.0003	11$\pm$2	50	2.2$\pm$0.3	-16.5$\pm$0.2	0.0017$\pm$0.0003	14$\pm$3
I4	20	3.2$\pm$0.8	-15.3$\pm$0.3	0.003$\pm$0.001	7$\pm$2	20	1.4$\pm$0.3	-16.4$\pm$0.1	0.006$\pm$0.001	5$\pm$1
	30	3.2$\pm$0.8	-15.3$\pm$0.3	0.002$\pm$0.001	9$\pm$2	30	1.4$\pm$0.3	-16.4$\pm$0.1	0.003$\pm$0.001	6$\pm$2
	40	3.2$\pm$0.8	-15.3$\pm$0.3	0.0014$\pm$0.0004	11$\pm$3	40	1.4$\pm$0.3	-16.4$\pm$0.1	0.003$\pm$0.001	8$\pm$2
	50	3.2$\pm$0.8	-15.3$\pm$0.3	0.0011$\pm$0.0003	14$\pm$3	50	1.4$\pm$0.3	-16.4$\pm$0.1	0.002$\pm$0.001	10$\pm$3

Table 9. Table 9: Values obtained with the fitting procedure described in Sect.  3.3 to the N 2 H + (1–0), 15 NNH + (1–0) and N 15 NH + (1–0) transitions towards D1, D2 and D3. The upper limits and tentative detections are obtained as explained in Sect.  3.3.1 .

Source	N2H+(1–0)
	$T_{\mathrm{ex}}$	$FWH{M}^a$	$v_{\mathrm{LSR}}$	$\tau$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1013cm-2)
D1	20	10.0$\pm$0.7	-15.0$\pm$0.4	0.14$\pm$0.02	4.9$\pm$0.8
	30	10.0$\pm$0.7	-15.0$\pm$0.4	0.087$\pm$0.008	6$\pm$1
	40	10.1$\pm$0.7	-15.0$\pm$0.4	0.063$\pm$0.006	8$\pm$2
	50	10.1$\pm$0.7	-15.0$\pm$0.4	0.049$\pm$0.004	10$\pm$2
D2	20	8.9$\pm$0.9	-14.4$\pm$0.4	0.12$\pm$0.02	3.7$\pm$0.7
	30	9.0$\pm$0.9	-14.4$\pm$0.4	0.075$\pm$0.009	4.9$\pm$0.9
	40	9.0$\pm$0.9	-14.4$\pm$0.4	0.054$\pm$0.007	6$\pm$1
	50	9.0$\pm$0.9	-14.4$\pm$0.4	0.042$\pm$0.005	7$\pm$2
D3	20	9$\pm$1	-14.8$\pm$0.5	0.11$\pm$0.02	3.4$\pm$0.7
	30	9$\pm$1	-14.7$\pm$0.5	0.069$\pm$0.01	4.5$\pm$0.9
	40	9$\pm$1	-14.7$\pm$0.5	0.050$\pm$0.007	6$\pm$1
	50	9$\pm$1	-14.7$\pm$0.5	0.039$\pm$0.006	7$\pm$2
Source	15NNH+(1–0)					N15NH+(1–0)
	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)
D1	20				$\le$2	20				$\le$2.4
	30				$\le$2.6	30				$\le$3.2
	40				$\le$3.3	40				$\le$4
	50				$\le$4	50				$\le$4.9
D2	20	0.8$\pm$0.2	-15.1$\pm$0.1	0.002$\pm$0.001	1.1$\pm$0.4c	20				$\le$2.4
	30	0.8$\pm$0.2	-15.1$\pm$0.1	0.0014$\pm$0.0005	1.5$\pm$0.6c	30				$\le$3.3
	40	0.8$\pm$0.2	-15.1$\pm$0.1	0.0010$\pm$0.0004	1.9$\pm$0.7c	40				$\le$4.2
	50	0.8$\pm$0.2	-15.1$\pm$0.1	0.0008$\pm$0.0003	2.4$\pm$0.8c	50				$\le$5
D3	20				$\le$1.4	20				$\le$1
	30				$\le$1.8	30				$\le$1.4
	40				$\le$2.3	40				$\le$1.7
	50				$\le$2.8	50				$\le$2.1

Table 10. Table 10: Values obtained with the fitting procedure described in Sect.  3.3 to the N 2 H + (1–0), 15 NNH + (1–0) and N 15 NH + (1–0) transitions towards mm1 and mm3 in regions equals to the IRAM-30m beam (see Sect  3.3.2 ).

Source	N2H+(1–0)
	$T_{\mathrm{ex}}$	$FWH{M}^a$	$v_{\mathrm{LSR}}$	$\tau$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1013cm-2)
mm1	43	10.0$\pm$0.7	-14.8$\pm$0.4	0.061$\pm$0.006	9$\pm$2
mm3	18	9.9$\pm$0.6	-14.7$\pm$0.3	0.18$\pm$0.02	5.0$\pm$0.8
Source	15NNH+(1–0)					N15NH+(1–0)
	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$	$T_{\mathrm{ex}}$	$FWHM$	$v_{\mathrm{LSR}}$	${\tau }^b$	$N_{\mathrm{tot}}$
	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)	(K)	(km s-1)	(km s-1)		($\times$1011cm-2)
mm1	43	1.7$\pm$0.3	-16.1$\pm$0.1	0.0009$\pm$0.0001	4.1$\pm$0.7	43	2.1$\pm$0.4	-16.2$\pm$0.2	0.0008$\pm$0.0002	5$\pm$2
mm3	18	1.9$\pm$0.4	-15.8$\pm$0.1	0.0022$\pm$0.0004	2.3$\pm$0.4	18	1.8$\pm$0.3	-16.1$\pm$0.2	0.003$\pm$0.001	2.6$\pm$0.7

Equations

$${\mathrm{N}}_2{\mathrm{H}}^{+}+{}^{15}\mathrm{N}\leftrightarrow {\mathrm{N}}^{15}{\mathrm{NH}}^{+}+\mathrm{N},$$

$${\mathrm{N}}_2{\mathrm{H}}^{+}+{}^{15}\mathrm{N}\leftrightarrow {}^{15}\mathrm{N}{\mathrm{NH}}^{+}+\mathrm{N}.$$

Full text

††pubyear: 2018††pagerange: First interferometric study of enhanced N-fractionation in N2H*+*: the high-mass star-forming region IRAS 05358+3543–LABEL:lastpage

1 Introduction

The origin of the Solar System is still a highly debated topic. Understanding the origin of the molecules in different Solar System bodies is important to have information on the proto-Solar nebula (PSN), from which our Sun was born. In particular, it is not clear if molecules in pristine Solar System materials, like comets or asteroids, were inherited from the cold and dense PSN, or if they are the result of chemical processing within the Solar protoplanetary disc.

One of the most used approaches involves the measurement of the 14N/15N ratio from different molecules. The ratio measured for the PSN from the Solar wind is 441$\pm$6 (Marty et al. 2010), and this value is about two times larger than that measured in the terrestrial atmosphere (TA), derived from N2, $\sim$272 (Marty et al. 2009). Colzi et al. (2018b) derived a new 14N/15N ratio in the local ISM of 375$\pm$60, broadly consistent with the PSN value and consistent with the values predicted by the Galactic chemical evolution model of Romano et al. (2017). The PSN value is also larger than that measured in some comets ($\sim$140, observed in CN, HCN and NH2, average value from Hily-Blant et al. 2017 and reference therein) and in carbonaceous chondrites ($\sim$50-250, e.g. van Kooten et al. 2017). van Kooten et al. (2017) found also a different 15N-enrichment between the amine groups and the nitrile functional groups in the lithic clasts of the Isheyevo carbonaceous chondrite of their study. New measurements also confirm the trend of low 14N/15N values in pristine Solar System materials, in fact Yang et al. (2017) derived a 14N/15N=130$\pm$15 from CN and 14N/15N=140$\pm$28 from NH2 in the outbursting comet C/2015 ER61. Moreover, McGeoch et al. (2018) determined, with a study of polymers of amino acids, in Allende and Acfer 086 meteorites 14N/15N of $\sim$200 and $\sim$140, respectively.

The relation between the 15N-enrichments in pristine Solar System material and the natal core is still uncertain, and probably these differences are also coming from different nitrogen reservoirs. 14N/15N ratios were measured through observations towards both low- and high-mass star-forming regions from different molecules: HCN and HNC (e.g. Hily-Blant et al. 2013a, Zeng et al. 2017, Colzi et al. 2018a, 2018b, Magalhães et al. 2018), N2H*+(e.g. Daniel et al. 2016) and NH3* (e.g. Gerin et al. 2009). None of these values are consistent with the anti-fractionation (14N/15N>500) observed with N2H*+* by Bizzocchi et al. (2013), Redaelli et al. (2018) and Fontani et al. (2015). Moreover, a difference between nitrile molecules (i.e. CN-bearing, such as HCN/HNC) and hydrogenated N-bearing molecules (e.g. NH3 and N2H*+) has been claimed, e.g. by Hily-Blant et al. (2013a) and Wampfler et al. (2014). A different behaviour between HCN/HNC and N2H+* is also apparent in massive star-forming regions. In fact, Fontani et al. (2015) found a larger dispersion of values from N2H*+*(from 180 up to 1300), with respect to the values found by Colzi et al. (2018a) from HCN and HNC (from 250 up to 650) in the same sample of sources in which were also performed other astrochemical studies (e.g. Colzi et al. 2018a, Mininni et al. 2018 and Fontani et al. 2018).

A summary of the measured 14N/15N ratios from different molecules can be found in Wirström et al. (2016).

Guzmán et al. (2017) have measured the 14N/15N from HCN towards 6 protoplanetary discs, and found that most of them shows 14N/15N consistent with comet-like values (80-160, with an average value of 111$\pm$9). However, Hily-Blant et al. (2017) measured a CN/C15N ratio of 323$\pm$30 in the disk orbiting the nearby young star TW Hya, and they proposed the CN as a possible present-day reservoir of nitrogen in the Solar neighbourhood.

From a theoretical point of view, two possible mechanisms were proposed to explain 15N-fractionation: isotope-exchange reactions, favoured at low temperatures (T<20K, e.g. Roueff et al. 2015, Wirström & Charnley 2018, Loison et al. 2019), or selective photo-dissociation of 14N15N over 14N2 (Heays et al. 2014), favoured in the external layers of the discs exposed directly to the illumination of the central (proto-)star. Visser et al. (2018) confirmed that isotope-selective photodissociation could be crucial to understand the N-fractionation in protoplanetary discs, especially those exposed to strong irradiation fields. In particular, for N2H*+*, this is true if the dominant fractionation mechanism are reactions with atomic N.

In this work we will focus on N2H*+. Terzieva & Herbst (2000) proposed that the reactions that cause most of 15N-enrichment in N2H+*are:

[TABLE]

However, the most recent chemical models have challenged this scenario, due to the recent discovery that reactions (1a) and (1b) do not occur in cold environments due to the presence of an entrance barrier (Roueff et al. 2015, Wirström & Charnley 2018). These chemical models fail to reproduce both the anti-fractionation (i.e. 14N/15N>500) measured in cold prestellar cores (Redaelli et al. 2018) and the large spread of values measured from different high-mass star forming regions (Fontani et al. 2015). This also suggests that chemical reactions may still be missing in existing models, or that the 15N-enrichment (or the contrary, i.e. 15N-depletion, which causes anti-fractionation) is a local phenomenon occurring in spatial regions smaller than the beam size of single-dish observations described above. In fact, for example with the single dish observations made by Fontani et al. (2015), the isotopic ratios obtained were average values over angular sizes of $\sim$30”, which hence include both the large-scale envelope and the dense core(s) embedded within it. A more recent chemical model of Furuya & Aikawa (2018) proposed that anti-fractionation could be explained introducing other mechanisms in chemical models, like isotope selective photodissociation of N2 and grain surface chemistry. We will discuss this point in Sect. 4.2.

In this work we report, for the first time, an interferometric analysis of the isotopic ratio 14N/15N towards the high-mass star-forming protocluster IRAS 05358+3543 (hereafter 05358) from N2H*+. We discuss the possible different N-fractionation between 15NNH+* and N15NH*+* and the possible 15N-enrichment at core scales ($\sim$5”). N-fractionation in N2H*+*towards this source was already studied by Fontani et al. (2015), who analysed IRAM-30m observations towards a sample of 26 high-mass star-forming regions, including 05358. In particular, for this source they found that the 14N/15N is lower ($\sim$200) than in the other massive star-forming regions.

The source and the observations are described in Sect. 2, the results are presented in Sect. 3 and a discussion of the results is presented in Sect. 4.

2 Source and observations

05358 is part of a sample of 69 high-mass protostellar objects studied by Sridharan et al. (2002), Beuther et al. (2002b), (2002c) and (2002d) in the millimeter wavelengths. The source lies at a distance of 1.8 kpc in the Auriga molecular cloud complex (Heyer et al. 1996), in the Perseus spiral arm of the Milky Way, and has a bolometric luminosity of 6300 L*⊙* (Snell et al. 1990). The source is part of a complex group of Hii regions: SH 235 is the brightest of four optical Hii regions and it is excited by the O9.5 star BD+351201 (Georgelin 1975). The source 05358 is associated with the Hii region SH 233 which is approximately 25$\arcmin$ west and 5$\arcmin$ south of SH 235.

As evidence for massive star-formation, the source is associated with maser emissions (e.g. Minier et al. 2000, Hu et al. 2016) and massive outflows (e.g. Beuther et al. 2002a, Ginsburg et al. 2009). Beuther et al. (2007) resolved 05358 in at least three continuum sub-sources:

• mm1: it has a typical spectrum of a young massive protostar where the central source has already started hydrogen burning. The source could be a B1 Zero-Age-Main-Sequence star, with a luminosity of $\sim$103.72 L*⊙* and a stellar mass of $\sim$13 M*⊙* (Lang 1992). From the vicinity of the source, a collimated outflow is ejected with a rate of 6$\times$10*-4* M*⊙*/yr (Beuther et al. 2002a), indicating that the protostar is still accreting gas. Moreover, the source is associated with a 8.3 GHz emission (VLA data) that is likely from a hypercompact Hii region. The main (sub)mm continuum source, which is at the center of two molecular outflows (Beuther et al. 2002a), is resolved in two separate (sub)mm continuum peaks, mm1a and mm1b;

-mm2: it consists of at least two sub-sources, mm2a and mm2b, and the position of mm2b shifts from 1.2 mm to 875 $\mu$m. There is also the presence of other peaks (mm2c and mm2d). The source mm2a is detected as a compact one and this could indicate that only this source is a star-forming region, while the others could be transient sources caused by the multiple outflow system in the region;

-mm3: for this source there are neither cm nor compact line detection (Leurini et al. 2007) and then it could be a very cold massive core in an early evolutionary stage. The grey-body function used to fit the SED is constrained by the high-frequency ($\sim$700 GHz) measurement by Beuther et al. (2007) to dust temperatures below 20 K.

Table LABEL:coordinates contains the source coordinates.

We carried on observations with the IRAM NOEMA Interferometer of the J=1–0 transition of N2H*+, 15NNH+* and N15NH*+* towards 05358. Rest frequencies are 93.1734 GHz (Cazzoli et al. 2012), 90.2638 GHz and 91.2057 GHz (Dore et al. 2009), for the J=1–0 transition of N2H*+, 15NNH+* and N15NH*+, respectively. In Table LABEL:frequencies the energy of the upper levels, the Einstein coefficients and rest frequencies of each hyperfine transition are given, except for N2H+(1–0) for which only the rotational frequency is given, since the hyperfine structure is not resolved in our spectra. To discuss our results, in Sect. 4.1 we will also use the averaged emission map of the J=5–4, K=0,1 transitions of CH3CN. Rest frequencies are 91.987 GHz and 91.985 GHz, respectively (Cazzoli & Puzzarini 2006). Because these lines were in the same spectral band of the J=1–0 transition of N2H+*, synthesized beam, spectral resolution and rms are the same.

Observations were carried out in 4 days from Sep. 29 to Nov. 3, 2016, with 8 antennas in C and D configurations, providing baselines in between 15 and 240 m, corresponding to an angular resolution of $\sim 3.1\times 3\arcsec$. The amount of precipitable water vapour was generally in between 5 and 10 mm. Visibility amplitudes and phases were calibrated on 0552+398 and 0548+378. The absolute flux density scale was calibrated on MWC349 and LKHA101. The bandpass calibration was performed on 3C84 or 3C454.3. To incorporate short spacings to the interferometric maps, we have reduced and analysed IRAM-30m observations of the J=1–0 transition of N2H*+, 15NNH+* and N15NH*+. These observations were carried out on March 5, 6 and 7, 2017. We have obtained large-scale maps on an angular region of $\sim 120\arcsec$, i.e. about twice the NOEMA primary beam at the frequency of the N2H+*(1–0) transition. The data were obtained with the on-the-fly mapping mode. Pointing was checked every hour on nearby quasars. Focus was checked at start of the observations, and after sunset. Creation of the synthetic visibilities and merging of the two data sets were performed through standard procedures available in the GILDAS software package MAPPING.

Calibration and imaging were performed using CLIC and MAPPING software of the GILDAS package***The GILDAS software is available at http://www.iram.fr/IRAMFR/GILDAS. The synthesized beam, the final velocity resolution and the rms (root mean square) of the datacubes are given in Table LABEL:beam. The analysis of the data was done with MADCUBA†††Madrid Data Cube Analysis on ImageJ is a software developed in the Center of Astrobiology (Madrid, INTA-CSIC) to visualise and analyse single spectra and datacubes (Martín et al., in prep., Rivilla et al. 2016, Rivilla et al. 2017). software package. For the analysis we have used the spectroscopic parameters from the CDMS molecular catalog‡‡‡https://www.astro.uni-koeln.de/cdms/ (Müller et al. 2001, 2005; Endres et al. 2016).

3 Results

3.1 Continuum map

In Fig. 1 the 3 mm continuum map towards the source is shown. With this map we are able to distinguish the already known continuum millimeter sources mm1, mm2 and mm3 (Beuther et al. 2007). However, we are not able to resolve mm1 as a binary system composed by mm1a and mm1b (Table LABEL:coordinates) because of the insufficient spatial resolution (Table LABEL:beam). Moreover, we have observed the continuum emission of a new millimeter source (mm4) at right ascension $\alpha_{\rm 2000}$= 05h:39m:12.46s and declination $\delta_{\rm 2000}$=35°:45$\arcmin$:40.9$\arcsec$. Since this source is out of the primary beam at 1 mm, it was not analysed by Beuther et al. (2007) that focused their work in characterising the properties of 05358 at the center of the field of view.

3.2 Morphology of N2H*+*and 15N-isotopologues emission

In Fig. 2 we show the averaged emission map of N2H*+(1–0), which arises mainly from 3 cores: one associated with both mm1 and mm2 not resolved, another associated with mm3, and finally one associated with mm4. We have also highlighted the presence of two other N2H+(1–0) emission regions (A and B) and we will discuss them in Sect. 4.1. It can also be noted that the N2H+*(1–0) emission peaks are shifted by $\sim$2–3$\arcsec$ with respect to the mm1, mm2 and mm3 continuum sources. We will discuss this result in Sect. 4.1.

In Fig. 3 the averaged emission maps of the 15NNH*+(1–0) and N15NH+* (1–0) transitions are shown. They are obtained by integrating the lines over the channels with signal, above the 3$\sigma$ level, in the Narrow spectra. In particular, for N15NH*+(1–0) we were able to resolve the hyperfine structure in three distinct components (see Table LABEL:frequencies). However, since the J=1–0, F=1-1 and J=1–0, F=0-1 transitions are near the noise level, we have decided to create the emission map integrating only over the channels that correspond to strongest component, namely the J=1-0, F=2-1 transition. Moreover, in Fig. 3 the blue contours correspond to the polygons from which we have extracted the spectra, in order to derive the total column densities of the three molecules, and the corresponding 14N/15N ratios, in eight different regions in 05358. The polygons are defined from the 5$\sigma$ of the corresponding averaged map. In particular, P1a, P2a, P3a and P4a correspond to a value of 5$\times$0.72 mJy/beam km s-1* (from 15NNH*+, left panel of Fig. 3), while P1b, P2b, P3b and P4b correspond to a value of 5$\times$0.5 mJy/beam km s-1* (from N15NH*+*, right panel of Fig. 3).

We have also extracted spectra from polygons defined as the intersection of the 5$\sigma$ of the two 15N-isotopologues averaged maps (I1, I2, I3 and I4, left panel of Fig. 4). This has been done in order to compare the 14N/15N ratios obtained from 15NNH*+* and N15NH*+* from the same regions of the source. Finally, we have decided to estimate the 14N/15N ratios in the regions defined in the right panel of Fig. 4, D1, D2 and D3, less bright and corresponding to diffuse N2H*+*emission.

3.3 Fitting procedure and column density calculation

To fit the lines and to compute the total column densities from the extracted spectra we have used MADCUBA assuming local thermal equilibrium (LTE) conditions. This is a reasonable assumption since the observed cores present volume densities of the order of 106 cm*-3* (see Beuther et al. 2007), that is higher than the critical densities of the observed lines ($\sim$1.5–2$\times$105 cm*-3*, assuming kinetic temperatures from 20 K up to 50 K).

The MADCUBA-AUTOFIT tool takes into account four parameters to create a synthetic LTE line profile: total column density (N), excitation temperature ($T_{\rm ex}$), peak velocity (v), and full width half maximum ($FWHM$). We have fitted all the lines fixing the $T_{\rm ex}$. We have performed the analysis with different $T_{\rm ex}$ values (20, 30, 40 and 50 K) since we had only one transition for each molecules and we were not able to derive it directly from the data. These temperature values are also consistent with the kinetic temperatures derived for 05358 from NH3(1,1) by Lu et al. (2014). Leaving free the other parameters ($N$, $v$, and $FWHM$) the AUTOFIT tool compares the LTE synthetic spectra with the observed spectra and provides the best non-linear least squared fit using the Levenberg-Marquardt algorithm. When the algorithm converges, it provides also the associated errors to the parameters. From these parameters, the tool calculates also the line opacity between levels $i$ and $j$, $\tau_{\rm ij}$ (see eq. (1) and (2) in Rivilla et al. 2019).

The three lines that we have observed, the N2H*+(1–0), 15NNH+(1–0) and N15NH+(1–0) transitions, have hyperfine structure. We can resolve it only for N15NH+(1–0) since the line width found ($\sim$2 km s-1*) is smaller than the separation in velocity of the hyperfine components (e.g. Fig. 5). However, we have fitted the spectra of both the 15NNH*+(1–0) and N15NH+(1–0) transitions taking the hyperfine structure into account, in order to be consistent with the analysis of the two different 15N-isotopologues of N2H+. Conversely, for N2H+(1–0) we have performed the analysis taking into account only the rotational transition J=1–0, as if it was a single line. We did this because the low spectral resolution ($\sim$6.5 km s-1*) permits the convergence of the fits only fixing a FWHM>6.5 km s*-1*, that is not the real one of the source ($\sim$2 km s*-1* from the fits of 15N-isotopologues). The fit to a single transition permits to leave the FWHM as a free parameter and, even if the final FWHM is not reliable, the convergence of the fit is more precise. In Appendix B.2 a detailed analysis demonstrates that the use of this method gives similar results of the analysis of the same N2H*+*simulated spectrum where the hyperfine structure could be resolved, within the errors.

All the extracted spectra from the P1a–P4a, P1b–P4b, I1, I2, I3 and I4 regions, and the more diffuse N2H*+*emission regions D1, D2 and D3 (see Sect. 3.2) are shown in Appendix A. Column densities, FWHM, velocities and opacities derived from the fit of the extracted spectra are given in Appendix C.

For the main isotopologue, N2H*+, we have found column densities of the order of 1014cm-2* in all regions. In particular, column densities towards P1a and P2a are higher than those towards P3a and P4a, of a factor $\sim$1.5. The same trend is followed by P1b and P2b with respect to P3b and P4b, with a decrease of a factor $\sim$3. However, for the 15N-isotopologues column densities are the same in each region, within the errors, and are of the order of 1011cm*-2*–1012cm*-2*, depending on the assumed Tex. In fact, as a general trend for all the analysed molecules, N2H*+, 15NNH+* and N15NH*+*, column densities increases with the Tex, but their ratios remain the same, within the uncertainties.

3.3.1 Diffuse regions

The 15NNH*+(1–0) and N15NH+(1–0) transitions are not detected in any of the diffuse N2H+emission regions D1, D2 and D3, except for 15NNH+(1–0) in D2 for which we have obtained a tentative detection (see Figs. 11). We have considered tentative detections the lines where the peak synthesized beam temperature ($T_{\rm SB}^{\rm peak}$) was between 2.5$\sigma$ and 3$\sigma$ ($2.5\sigma\leq T_{\rm SB}^{\rm peak}<3\sigma$) , where $\sigma$ is the r.m.s noise of the spectrum. For the non detections we have computed upper limits for the column densities with MADCUBA. They have been derived from the formula 3$\sigma\times\Delta$v/$\sqrt{N_{\rm chan}}$, where $N_{\rm chan}$ is the number of channels covered by the line width $\Delta$v. In particular, we have considered as $\Delta$v an average value of the FWHM derived from the same transitions towards P1a–P4a and P1b–P4b (Table LABEL:polygon-fit): $\Delta$v${}_{{}^{15}{\rm NNH}^{+}}=2.3\pm 0.2$ km s-1* and $\Delta$v${}_{{\rm N}^{15}{\rm NH}^{+}}=1.7\pm 0.4$ km s*-1*. The upper limits and tentative detection fit results from diffuse N2H*+*emission regions are reported in Table LABEL:diffuse-fit.

3.3.2 Extracted spectra from IRAM-30m beam

In order to compare our results with those obtained by Fontani et al. (2015) with IRAM-30m observations, we have also extracted spectra from regions that correspond to the IRAM-30m beam around the sources mm1 and mm3, as defined in Table LABEL:coordinates (considering mm1a and mm1b at the same coordinates). This was also done in order to show how much the 14N/15N ratios change going to smaller scales with respect to those observed with single-dish observations. The IRAM-30m beam at 93 GHz is $\sim$27.6$\arcsec$. In order to be consistent with Fontani et al. (2015), we have fitted our spectra fixing the $T_{\rm ex}$ equal to those used in their work: 43 K for mm1 and 18 K for mm3. We noted that even though the two main beams are partially overlapping, the dominant emission of each of the two is from the corresponding central source. The extracted spectra from regions equivalent to the IRAM-30m beam are shown in Fig. 12 and the results from the fit are listed in Table LABEL:30m-fit.

3.4 14N/15N ratios

We have computed the 14N/15N ratios for N2H*+*, along with the uncertainties derived propagating the error on total column densities. In Tables LABEL:ratio-table1, LABEL:ratio-table2 and LABEL:ratio-table3 the 14N/15N ratios for the different regions defined in Sect. 3.2 are given. Since the ratios are obtained from column density values obtained from observations at different spectral resolution, we have carefully evaluated the effects of this data inhomogeneity in Appendix B. There we show that the different spectral resolution does not affect significantly the column densities, and hence the 14N/15N ratios.

In general, we have found values that span from $\sim$100 up to $\sim$220, where the higher values are comparable with those observed with the IRAM-30m (Fontani et al. 2015). Moreover, the assumption of different $T_{\rm ex}$ does not change the results, that are consistent within the errors. Therefore, in the next session we will discuss the final results considering a single excitation temperature. Moreover, a different $T_{\rm ex}$ can change the column densities of each molecule, but does not influence their ratios, since the dependence on $T_{\rm ex}$ in the ratio is irrelevant when the two column densities are both computed from the (1–0) transition and the lines are not optically thick. This is supported by the fact that the ratios given in Table LABEL:ratio-table1 are the same within the uncertainties, towards each region.

4 Discussion

4.1 Comparison between line and continuum emission maps

The averaged emission maps of the J=1–0 transition of N2H*+, 15NNH+* and N15NH*+* (Fig. 2 and 3) present a shift of $\sim$2-3$\arcsec$ to the south-west with respect to the continuum sources. In particular this effect is stronger for mm1, mm2 and mm3 than for mm4. The overall structure of the integrated N2H*+(1–0) (and its 15N-isotopologues) consists of three main cores. Two of these are located to the south of the mm1 and mm2 continuum sources, and to the west of the mm3 continuum sources, while the third is located exactly to the mm4 source. Numerous studies have shown that during the collapse of a low-mass core, the CO desorbs from the grain mantles, because of the rising of the temperature, and this causes a substantial destruction of N2H+(e.g. Di Francesco et al. 2004). The chemical processes at work in high-mass star-forming regions are likely different from those in low-mass star-forming ones, but the destruction of N2H+* by CO should occur in both cases. Towards mm1 the presence of a young massive protostar is well-known (see Sect. 2) and the displacement between the dust and N2H*+peak suggests that the heating of the protostar, and probably of the molecular outflows associated with it (Beuther et al. 2002a), may have caused the desorption of CO from grain mantles and the subsequent destruction of N2H+. Probably something similar is happening also around mm3 although we do not have clear evidences that star formation already started. Moreover, the correspondence of the dust and N2H+peak towards mm4 is an evidence that it is chemically less evolved and hence maybe starless. Something similar for high-mass star-forming regions has been reported towards IRAS 23033+5951 (Reid & Matthews 2008) and towards the massive protocluster AFGL 5142 (Busquet et al. 2011). To investigate better this situation we have searched in our spectral setup for other molecules which could better trace the continuum sources and the associated star formation activity. In particular, in Fig. 6 the averaged emission map of CH3CN(J=5–4, K=0,1) is shown. In the left panel the clear correspondence between the 3mm-continuum sources and CH3CN can be noted; moreover, in the right panel it is evident that CH3CN and N2H+do not coincide towards mm1, mm2 and mm3, while the core-structure defined by the two different molecules corresponds towards mm4. Moreover, the asymmetry of N2H+*around the millimeter continuum sources could be due to the presence of the interaction with the gas of the complex group of Hii regions located in the North-East. However, we do not have observational proof of this point and it needs to be tested.

As already mentioned in Sect. 3.2, in the western and in the south-eastern part of the cluster, two other N2H*+-emission zones are present. In particular, A corresponds to a core structure in the 15N-isotopologues maps (P4a and P4b in Fig. 3) and corresponds also to one of the emission peaks of the high-density tracer H13CO+* defined in Fig. 8 of Beuther et al. (2002a), which is a high-density tracer. However, B presents a more complicated structure, probably related with the presence of multiple outflows in the whole region (outflow A in Beuther et al. 2002a). Leurini et al. (2011) discussed the N2H*+(1–0) emission towards the IRDC G351.77-0.51 and they found a velocity gradient on large scale, probably associated with outflows detected in CO. Although N2H+is seldom associated with outflows, this could be due to an in interaction of outflow shocks with the molecular envelope (traced also by N2H+). A similar situation was detected in the same transitions towards the class 0 object IRAM 04191+1522 by Lee et al. (2005). Moreover, Codella et al. (2013) present the first detection of N2H+*towards a low-mass protostellar outflow, L1157-B1, a bow-shock, at $\sim$0.1 pc from the protostellar cocoon L1157.

4.2 Is N-fractionation a core-scale effect?

Fontani et al. (2015) have derived, with the IRAM-30m radiotelescope, 14N/15N ratios of 190$\pm$20 with 15NNH*+* and of 180$\pm$23 with N15NH*+* towards mm1, and of 210$\pm$12 with 15NNH*+* and of 180$\pm$13 with N15NH*+* towards mm3. Extracting the spectra in regions equivalent to the IRAM-30m beam at 93 GHz, we have derived the 14N/15N ratios listed in the lower panel of Table LABEL:ratio-table3. These are consistent with the results obtained with the single-dish telescope, within the errors. The comparison of these results with the 14N/15N ratios obtained towards P1a, P2a, P3a and P4a and towards P1b, P2b, P3b and P4b is shown in Fig. 7. In particular, it is evident from left panel of Fig. 7 that there is a 15N-enrichment (i.e. a 14N/15N ratio significantly lower) towards P2a, P3a and P4a. The same happens towards P2b and P3b (right panel of Fig. 7). This means that at higher angular resolution, where the core scales can be resolved, it is possible to reveal a different N-fractionation with respect to that observed with a single-dish telescope and more representative of the average 14N/15N between the cores and the envelope of the whole star-forming region. This suggests that the mechanisms that produce more 15N with respect to 14N in N2H*+*can occur in smaller regions and could be a local effect. The fact that towards P1a and P1b, 14N/15N values consistent with IRAM-30m observations ($\sim$200) are found could be explained by the fact that these regions are near an evolved high-mass source (mm1) and the chemistry could be different toward hotter gas. Moreover, the different 14N/15N ratios between P4a and P4b, $\sim$115 and $\sim$185, respectively, could be due to a different behaviour of the two 15N-isotopologues, but we will discuss about this in the next section.

In general, the lower 14N/15N ratios of $\sim$100–110, are similar to the values measured by Guzmán et al. (2017) toward protoplanetary discs in HCN, and they are also consistent with the low values measured in pristine Solar System materials. However, N-fractionation of different molecules, like HCN and N2H*+*, could be different. To investigate better this point, consistent measurements of N-fractionation from different species in the same sources are needed to understand the puzzle of nitrogen isotopic ratios. Moreover, high-mass angular resolution observations towards other high-mass star-forming regions are needed in order to confirm these results. In fact, it is also important to gather more data in sources that are good candidates to represent the environment in which our Sun was born. In this respect, intermediate- and high-mass star-forming cores are interesting targets because the Sun was probably born in a rich cluster that also contained massive stars (e.g. Adams 2010, Lichtenberg et al. 2019).

Unfortunately, none of the published chemical models up to now are able to reproduce the low 14N/15N observed in this work. Moreover, they are not able to reproduce observed values of N-fractionation in N2H*+in general. However, they are gas-phase models that feature only low-temperature exchange reactions as the only fractionation mechanism (e.g. Roueff et al. 2015, Wirström & Charnley 2018). Probably other mechanisms must be taken into account, like the isotope selective photodissociation of N2* and grain surface chemistry (Furuya & Aikawa et al. 2018). In particular, these authors have proposed some mechanisms to explain the 15N-depletion in N2H*+in prestellar cores starting from their parental clouds. They found that during the evolution of a molecular cloud, where the external UV radiation field is not fully shielded, the 14N15N molecule could be photodissociated where the 14N2* is not, causing an enrichment of atomic 15N. The subsequent adsorption of atomic N creates an enrichment of 15N in NH3 ice by surface reactions. They found that as long as the photodesorption of NH3 ice is less efficient than that of N2 ice, the net effect is the loss of 15N from the gas phase. However, these results are based on a specific physical model adapted for cloud with density of $\sim$104cm*-3*, which could not be appropriate for dense sources. To understand the existing 15N-observations more numerical studies of 15N-fractionation are needed.

4.3 15NNH*+* vs N15NH*+*

In this section we will discuss about the possible differences between the two 15N-isotopologues of N2H*+: 15NNH+* and N15NH*+. In order to do that, we have compared the 14N/15N ratios obtained from 15NNH+* and N15NH*+* towards I1, I2, I3 and I4, since the column densities of the two 15N-isotopologues were derived from the same regions.

In Fig. 8 the 15NNH*+/N15NH+* column density ratios, for the different excitation temperatures assumed, are shown. In this case, the dependence on the $T_{\rm ex}$ is small because both the 15N-isotopologues are almost optically thin. We have found that 15NNH*+* is more abundant (at 1$\sigma$ level) with respect to N15NH*+* towards I4, and the contrary towards I1, I2 and I3. Concerning what happens towards I4, it is difficult to interpret our results from an evolutionary point of view, because the nature of this source is unknown (as already mentioned in Sect. 4.1). In this case, the ratio between the two 15N-isotopologues is higher than one, also including the errorbar. This could be a clue that in some particular cases of densities and temperatures the two 15N-isotopologues could form through different reaction pathways. This finding should also be confirmed by a higher statistics.

As already mentioned, the most recent chemical models are not able to reproduce a different behaviour between the two 15N-isotopologues. Other works toward very different environments show a tentative trend of 15NNH*+* to be less abundant than N15NH*+: Kahane et al. (2018) towards OMC-2 FIR4 (15NNH+/N15NH+$\sim$0.8) and Redaelli et al. (2018) towards the prestellar cores L1544 and L694-2 (15NNH+/N15NH+*$\sim$0.9 and $\sim$0.8, respectively).

4.4 14N/15N ratios in diffuse regions

In Tab. LABEL:ratio-table3 the results obtained towards the more diffuse N2H*+emission regions D1, D2 and D3 are shown. All of these results are lower limits since the 15N-isotopologues were not detected, except for 15NNH+* towards D2. The lower limit values obtained are, on average, higher than 200–250. These are also the maximum values obtained towards the peak-emission regions. Only the lower limits towards D2 from N15NH*+* are smaller, but given that the detection from 15NNH*+* confirm a 14N/15N ratio of $\sim$320, and that we have found a maximum difference of $\sim$1.5 (see previous Section), we can confirm that this lower limits corresponds to 14N/15N ratios higher than 200. These results mean that 15N tend to be less enriched in N2H*+in the outer parts of a star-forming region. This could be a possible confirmation of the work of Furuya et al. (2018): in NI/N2* transition zones, where the gas is less dense and the interstellar UV radiation field is not significantly attenuated ($A_{\rm v}$<3 mag), the 15N is frozen in NH3 ice and there is a loss of 15N from the gas phase.

5 Conclusion

We used the IRAM NOEMA Interferometer to observe the emission of N2H*+(1–0) and its 15N-isotopologues towards the high-mass star-forming protocluster IRAS 05358+3543. The N2H+dense gas emission consists of three main cores, two of which are displaced with respect to the 3mm continuum sources, probably because of the star-formation activity that causes the N2H+destruction by CO desorption from grain mantles. We have also discussed the presence of two other N2H+-emission zones: A and B. The first define a fourth core from the 15N-isotopologues emission maps and corresponds to a H13CO+* peak observed in previous works. The second is probably associated with an outflow coming out from the more evolved region in the north.

The 14N/15N ratios derived toward regions that correspond to the IRAM-30m beam are higher than those from the cores of the cluster. This indicates the possibility that 15N-enrichment in star-forming regions is a local effect. The lower values of $\sim$100 toward the cores are similar to some values measured in protoplanetary discs, and also to the low values measured in the pristine Solar System material. Moreover, the 14N/15N ratios derived from 15NNH*+* and N15NH*+* show some differences. In particular, towards I1 and I4 15NNH*+* is more abundant than N15NH*+. However, the most recent chemical models are not able to reproduce the low 14N/15N values found for N2H+*in this work and in other works.

Finally, the 14N/15N upper limits derived in the more diffuse N2H*+emission regions of the cluster, point to values higher than those derived in the cores. This confirms the conclusion that N-fractionation, measured from N2H+*, seems to be a local effect, at least in this high-mass star-forming region. However, more observations, also of different molecules, are needed to unveil which physical parameters are affecting the N-fractionation, thus helping us to gain understanding on the chemistry involved.

Acknowledgements

This work is based on observations carried out under projects number S16AK and 081-16 with the IRAM NOEMA Interferometer and 30m telescope, respectively. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain). We are grateful to the IRAM NOEMA staff for their help during the calibration of the data and the IRAM-30m 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. LC acknowledges support from the Italian Ministero dell’Istruzione, Università e Ricerca through the grant Progetti Premiali 2012 - iALMA (CUP C52I13000140001). PC acknowledges support from the European Research Council (project PALs 320620). LC sincerely acknowledge the hospitality of the Joint ALMA Observatory (JAO), where significant part of the manuscript was written.

Appendix A Extracted spectra

In this appendix, the extracted spectra of the J=1–0 transition of N2H*+, 15NNH+* and N15NH*+* from the regions P1a–P4a, P1b–P4b, I1, I2, I3, I4, D1, D2 and D3 are shown. The definition of the regions can be found in Sect. 3.

Appendix B Spectra simulation tests

In this appendix we have estimated how much our results are influenced by the different spectral resolution between the main isotopologue (N2H*+, of about 6.45 km/s), and the 15N-isotopologues (15NNH+* and N15NH*+*, of about 0.5 km/s). For this test we have taken into account the spectra extracted from P1a and P1b, and the corresponding fit results with a $T_{\rm ex}$=20 K

B.1 15N-isotopologues at lower spectral resolution

We have smoothed the 15NNH*+(1–0) transition to a 6.5 km/s resolution. In Fig. 13, the black spectrum is the observed one, the blue solid line is the fit to the observed spectrum, and the red line is the smoothed one. The integrated intensity (A=$\int T_{\rm MB}\,dv$) of the observed spectrum is (0.27$\pm$0.04) K km/s (relative error $\Delta A$/A$\simeq$15%), while the integrated intensity of the smoothed spectrum is (0.31$\pm$0.06) K km/s, 13% higher than the observed one. This error is comparable to the error given from the fit. We have done the same analysis with N15NH+*(1–0) (Fig. 14), for which we have obtained, for the smoothed spectrum, an integrated intensity of (0.26$\pm$0.08) K km/s, 12% higher than the observed one (0.23$\pm$0.03 K km/s), that again is comparable with the relative error given by the fit of the observed spectrum ($\Delta A$/*A$\simeq$*13%). The test hence shows that the integrated intensity of the 15N-isotopologues, and thus their column densities, are not affected by changing the resolution from 0.5 to 6.5 km/s.

B.2 N2H*+*at higher spectral resolution

We have simulated the spectrum of N2H*+(1–0) with a resolution of 0.5 km/s using the results of the fit obtained for a $T_{\rm ex}$=20 K, towards P1a, but with a line width of 2.3 km/s (which is the average value measured from 15NNH+). The obtained spectrum is the blue in Fig. 15, with an integrated intensity of (44$\pm$2) K km/s. Then, we have smoothed this simulated spectrum to 6.45 km/s obtaining the red one in the same figure. The integrated intensity of this latter is the same of that at higher spectral resolution. We can also note that the smoothed simulated spectrum is comparable with the observed one (black) that has an integrated intensity of (48$\pm$2) K km/s ($\Delta A$/A$\simeq$4%), 8% higher than the simulated one. We can conclude that, even if we cannot resolve the hyperfine structure of N2H+*(1–0), with our analysis we have obtained an accurate integrated intensity, with at most 4% overestimation.

Appendix C Fit results

In this appendix, the results of the fitting procedure to the N2H*+(1–0), 15NNH+(1–0) and N15NH+*(1–0) transitions towards all the sources defined in Sect. 3. The method is explained in Sect. 3.3.

a Note that these FWHM overestimate those obtained with the 15N-isotopologues since N2H*+*has a spectral resolution higher than the physical FWHM;

b opacity of the main hyperfine component.

a Note that these FWHM overestimate those obtained with the 15N-isotopologues since N2H*+*has a spectral resolution higher than the physical FWHM;

b opacity of the main hyperfine component.

a Note that these FWHM overestimate those obtained with the 15N-isotopologues since N2H*+*has a spectral resolution higher than the physical FWHM;

b opacity of the main hyperfine component;

c tentative detection.

a Note that these FWHM overestimate those obtained with the 15N-isotopologues since N2H*+*has a spectral resolution higher than the physical FWHM;

b opacity of the main hyperfine component;