Optical/NIR stellar absorption and emission-line indices from luminous infrared galaxies

TL;DR

This study provides a comprehensive set of optical and near-infrared stellar absorption and emission-line indices for luminous infrared galaxies, highlighting correlations and model comparisons to improve stellar population understanding.

Contribution

It offers a homogeneous dataset of 45 spectral indices for 16 infrared-luminous and 19 early-type galaxies, facilitating testing of stellar population models across optical and NIR wavelengths.

Findings

Optical indices agree well with models.

NIR indices are challenging to interpret.

Models match early-type spirals but not star-forming galaxies.

Abstract

We analyze a set of optical-to-near-infrared long-slit nuclear spectra of 16 infrared-luminous spiral galaxies. All of the studied sources present H$_2$ emission, which reflects the star-forming nature of our sample, and they clearly display H I emission lines in the optical. Their continua contain many strong stellar absorption lines, with the most common features due to Ca I, Ca II, Fe I, Na I, Mg I, in addition to prominent absorption bands of TiO, VO, ZrO, CN and CO. We report a homogeneous set of equivalent width (EW) measurements for 45 indices, from optical to NIR species for the 16 star-forming galaxies as well as for 19 early type galaxies where we collected the data from the literature. This selected set of emission and absorption-feature measurements can be used to test predictions of the forthcoming generations of stellar population models. We find correlations among the different absorption features and propose here correlations between optical and NIR indices, as well as among different NIR indices, and compare them with model predictions. While for the optical absorption features the models consistently agree with the observations,the NIR indices are much harder to interpret. For early-type spirals the measurements agree roughly with the models, while for star-forming objects they fail to predict the strengths of these indices.

Figures

Tables

Table 1. Table 1: Near-Infrared observation log and basic sample properties.

Source	$\alpha$	$\delta$	z	Obs. Date	Exp. Time	Airmass	PA	Size	Activity	log($\frac{L_{IR}}{L_{\odot }}){}^{\star }$	Morphology
					(s)		(deg)	(pc$\times$pc)		(13)
NGC 23	00h09m53.4s	+25d55m26s	0.0157202	2010 10 07	29 $\times$ 120	1.04	330	1348 $\times$ 270	SFG1	11.05	SBa
NGC 520	01h24m35.1s	+03d47m33s	0.0080367	2010 10 04	16 $\times$ 120	1.04	300	689 $\times$ 138	SFG2	10.91	S0
NGC 660	01h43m02.4s	+13d38m42s	0.0029152	2010 10 06	24 $\times$ 120	1.01	33	237 $\times$ 50	Sy2/HII2,3,4	10.49	SBa pec
NGC 1055	02h41m45.2s	+00d26m35s	0.0036267	2010 10 04	16 $\times$ 120	1.07	285	466 $\times$ 62	LINER/HII2,3,4	10.09	Sbc
NGC 1134	02h53m41.3s	+13d00m51s	0.0129803	2010 10 04	16 $\times$ 120	1.11	0	1113 $\times$ 223	SFG5	10.83	S?
NGC 1204	03h04m39.9s	-12d20m29s	0.0154058	2010 10 07	16 $\times$ 120	1.23	66	1321 $\times$ 264	LINER6	10.88	S0/a
NGC 1222	03h08m56.7s	-02d57m19s	0.0082097	2010 10 06	24 $\times$ 120	1.13	315	598 $\times$ 141	SFG7	10.60	S0 pec
NGC 1266	03h16m00.7s	-02d25m38s	0.0077032	2010 10 07	18 $\times$ 120	1.09	0	661 $\times$ 132	LINER7	10.46	SB0 pec
UGC 2982	04h12m22.4s	+05d32m51s	0.0177955	2010 10 04	9 $\times$ 120	1.11	295	1526 $\times$ 305	SFG8	11.30	SB
NGC 1797	05h07m44.9s	-08d01m09s	0.0154111	2010 10 07	16 $\times$ 120	1.23	66	1321 $\times$ 264	SFG1	11.00	SBa
NGC 6814	19h42m40.6s	-10d19m25s	0.0056730	2010 10 07	16 $\times$ 120	1.17	0	486 $\times$ 97	Sy 17	10.25	SBbc
NGC 6835	19h54m32.9s	-12d34m03s	0.0057248	2010 10 06	22 $\times$ 120	1.21	70	368 $\times$ 98	SFG9	10.32	SBa
UGC 12150	22h41m12.2s	+34d14m57s	0.0214590	2010 10 04	15 $\times$ 120	1.08	37	1656 $\times$ 368	LINER/HII10	11.29	SB0/a
NGC 7465	23h02m01.0s	+15d57m53s	0.0066328	2010 10 06	12 $\times$ 120	1.03	340	569 $\times$ 114	LINER/Sy 211	10.10	SB0
NGC 7591	23h18m16.3s	+06d35m09s	0.0165841	2010 10 07	16 $\times$ 120	1.03	0	1422 $\times$ 284	LINER7	11.05	SBbc
NGC 7678	23h28m27.9s	+22d25m16s	0.0120136	2010 10 04	16 $\times$ 120	1.01	90	927 $\times$ 206	SFG12	10.77	SBc

Table 2. Table 2: Optical observation log. The slit was oriented North-South.

Source	Obs. Date	Exp.	Airmass	Size
		Time (s)		(pc$\times$pc)
NGC 23	2010 10 04	600	1.20	2359 $\times$ 1348
NGC 520	2010 10 03	600	1.28	4307 $\times$ 689
NGC 660	2010 10 04	600	1.18	437 $\times$ 250
NGC 1055	2010 10 04	600	1.36	544 $\times$ 311
NGC 1134	2010 10 04	600	1.14	12243 $\times$ 1113
NGC 1204	2010 10 04	600	1.68	2312 $\times$ 1321
NGC 1222	2010 10 03	600	1.42	3344 $\times$ 704
NGC 1266	2010 10 03	600	1.39	6440 $\times$ 661
UGC 2982	2010 10 04	600	1.27	2670 $\times$ 1526
NGC 1797	2010 10 02	600	1.56	11893 $\times$ 1321
UGC 12150	2010 10 03	600	1.03	18401 $\times$ 1840
NGC 7465	2010 10 03	600	1.12	3270 $\times$ 569
NGC 7591	2010 10 03	600	1.21	13154 $\times$ 1422
NGC 7678	2010 10 04	600	1.18	1803 $\times$ 1030

Table 3. Table 3: Emission line fluxes in units of 1 ​ x ​ 10 − 15 ​ erg ​ cm − 2 ​ s − 1 1 x superscript 10 15 erg superscript cm 2 superscript s 1 \rm 1x10^{-15}\,erg\,cm^{-2}\,s^{-1} .

Line	Ion	NGC23	NGC520	NGC660	NGC1055	NGC1134	NGC1204	NGC1222	NGC1266
	$C_{ext}$	–	–	3.42$\pm$0.13	–	–	2.95$\pm$0.14	1.06$\pm$0.05	6.74$\pm$0.73
4861	H$\beta$	20.40$\pm$1.16	–	0.86$\pm$0.26	–	6.06$\pm$0.37	0.30$\pm$0.06	49.30$\pm$0.65	–
4959	[O iii]	30.70$\pm$3.13	–	3.37$\pm$0.41	–	8.25$\pm$0.72	1.69$\pm$0.48	47.30$\pm$0.75	–
5007	[O iii]	10.50$\pm$3.13	2.16$\pm$0.30	1.15$\pm$0.41	–	2.81$\pm$0.72	0.58$\pm$0.48	140.00$\pm$0.75	4.27$\pm$0.64
6548	[N ii]	108.00$\pm$3.26	1.60$\pm$0.31	6.25$\pm$0.53	1.41$\pm$0.32	12.90$\pm$1.41	5.75$\pm$0.43	14.20$\pm$0.40	14.30$\pm$1.07
6563	H$\alpha$	83.80$\pm$3.89	8.68$\pm$0.44	17.20$\pm$0.46	4.97$\pm$0.34	40.50$\pm$0.95	20.50$\pm$0.42	222.00$\pm$0.43	11.40$\pm$0.76
6583	[N ii]	27.90$\pm$3.89	6.96$\pm$0.49	18.60$\pm$0.48	4.56$\pm$0.39	38.30$\pm$1.21	18.60$\pm$0.42	49.70$\pm$0.45	30.60$\pm$0.87
6716	[S ii]	46.50$\pm$1.92	3.29$\pm$0.46	5.15$\pm$0.48	1.25$\pm$0.29	14.60$\pm$1.23	5.76$\pm$1.23	20.70$\pm$0.48	18.70$\pm$1.16
6730	[S ii]	34.70$\pm$1.92	3.07$\pm$0.55	4.88$\pm$0.61	1.31$\pm$0.44	12.90$\pm$1.50	5.69$\pm$1.23	19.90$\pm$0.57	21.10$\pm$1.16
9069	[S iii]	14.80$\pm$3.96	–	–	–	–	10.30$\pm$1.40	32.80$\pm$1.59	–
9531	[S iii]	15.70$\pm$3.96	–	23.90$\pm$0.94	–	–	11.90$\pm$0.54	72.60$\pm$1.49	–
9824	[C i]	1.57$\pm$0.60	–	2.10$\pm$0.27	–	–	1.15$\pm$0.30	1.55$\pm$0.49	–
9850	[C i]	5.15$\pm$0.60	–	2.53$\pm$0.27	–	–	1.88$\pm$0.30	0.97$\pm$0.49	9.21$\pm$1.06
10049	Pa$\delta$	–	–	2.55$\pm$0.22	–	–	1.03$\pm$0.08	4.05$\pm$0.31	–
10122	He ii	–	–	2.78$\pm$0.22	–	–	1.34$\pm$0.08	–	–
10830	He i	22.30$\pm$2.60	–	13.30$\pm$0.41	–	–	9.24$\pm$0.56	53.80$\pm$0.77	6.26$\pm$0.83
10938	Pa$\gamma$	5.65$\pm$1.58	–	8.57$\pm$0.31	–	–	2.78$\pm$0.27	12.90$\pm$0.76	–
11470	[P ii]	–	–	1.82$\pm$1.09	–	–	1.17$\pm$0.27	–	–
11886	[P ii]	–	–	3.68$\pm$1.09	–	–	1.70$\pm$0.19	–	–
12567	[Fe ii]	10.50$\pm$0.87	–	13.90$\pm$0.65	–	–	5.16$\pm$0.20	6.32$\pm$0.43	3.19$\pm$0.51
12820	Pa$\beta$	–	–	29.00$\pm$0.60	–	–	12.10$\pm$0.20	28.20$\pm$0.37	0.75$\pm$0.08
12950	[Fe ii]	–	–	1.21$\pm$0.15	–	–	1.35$\pm$0.25	–	–
13209	[Fe ii]	–	–	6.60$\pm$0.24	–	–	4.49$\pm$0.39	–	–
15342	[Fe ii]	–	–	–	–	–	1.52$\pm$0.36	–	–
16436	[Fe ii]	14.20$\pm$3.46	3.87$\pm$0.14	15.70$\pm$0.90	–	–	6.51$\pm$0.36	4.65$\pm$0.18	2.90$\pm$0.38
16773	[Fe ii]+Br11	31.50$\pm$1.34	–	–	–	–	2.65$\pm$0.50	–	–
17360	Br10	–	–	–	–	–	1.60$\pm$0.13	–	–
18750	Pa$\alpha$	–	35.40$\pm$0.30	60.20$\pm$1.64	–	–	61.10$\pm$0.33	69.30$\pm$0.57	8.03$\pm$0.34
19446	Br$\delta$	–	2.07$\pm$0.28	5.67$\pm$0.53	–	–	1.80$\pm$0.34	3.78$\pm$0.20	–
19570	H2	20.20$\pm$4.60	2.41$\pm$0.39	8.87$\pm$0.80	1.24$\pm$0.3	–	4.17$\pm$0.51	1.95$\pm$0.34	15.10$\pm$0.47
20332	H2	4.94$\pm$0.57	1.13$\pm$0.10	3.67$\pm$0.56	–	1.14$\pm$0.3	1.81$\pm$0.34	0.95$\pm$0.14	5.10$\pm$0.44
20580	H2	–	2.47$\pm$0.09	5.32$\pm$0.64	–	–	1.47$\pm$0.27	4.68$\pm$0.14	–
21218	H2	10.00$\pm$1.26	2.40$\pm$0.18	6.91$\pm$0.72	0.6$\pm$0.08	2.37$\pm$0.4	3.61$\pm$0.25	0.84$\pm$0.12	13.70$\pm$0.25
21654	Br$\gamma$	–	6.67$\pm$0.19	16.60$\pm$0.71	–	–	5.88$\pm$0.27	6.98$\pm$0.07	1.40$\pm$0.33
22230	H2	4.53$\pm$2.97	0.67$\pm$0.19	2.01$\pm$1.20	–	–	0.87$\pm$0.09	0.49$\pm$0.19	3.26$\pm$0.12
22470	H2	1.47$\pm$0.37	0.83$\pm$0.12	1.18$\pm$0.12	–	–	0.72$\pm$0.10	0.29$\pm$0.04	1.51$\pm$0.14

Table 4. Table 4: Continuation of Table  3

Line	Ion	UGC2982	NGC1797	NGC6814*	NGC6835	UGC12150	NGC7465	NGC7591	NGC7678
	$C_{ext}$	–	2.68$\pm$0.08	0.00	4.87$\pm$0.19	2.16$\pm$0.17	2.35$\pm$0.93	2.53$\pm$0.07	1.39$\pm$0.23
4861	H$\beta$	0.49$\pm$0.15	13.10$\pm$0.71	–	–	–	21.80$\pm$1.05	–	9.45$\pm$0.42
4959	[O iii]	0.35$\pm$0.16	–	–	–	–	31.40$\pm$2.27	–	–
5007	[O iii]	0.76$\pm$0.29	3.11$\pm$0.44	–	–	–	56.00$\pm$1.85	–	2.83$\pm$0.54
6548	[N ii]	1.64$\pm$0.39	11.20$\pm$0.59	–	–	–	29.50$\pm$1.03	8.15$\pm$1.36	7.88$\pm$0.82
6563	H$\alpha$	13.90$\pm$0.40	88.70$\pm$0.74	–	–	–	122.00$\pm$0.89	21.10$\pm$1.00	53.40$\pm$1.00
6583	[N ii]	5.62$\pm$0.39	46.40$\pm$0.76	–	–	–	71.10$\pm$0.98	18.60$\pm$1.04	27.60$\pm$1.00
6716	[S ii]	2.77$\pm$0.45	10.80$\pm$0.41	–	–	–	39.70$\pm$1.17	4.87$\pm$1.94	9.23$\pm$0.97
6730	[S ii]	2.66$\pm$0.75	10.30$\pm$0.46	–	–	–	34.10$\pm$1.22	3.46$\pm$1.94	10.30$\pm$1.37
9069	[S iii]	–	5.80$\pm$0.59	23.70$\pm$0.49	–	5.16$\pm$0.27	18.00$\pm$1.43	–	6.59$\pm$0.64
9531	[S iii]	–	12.10$\pm$0.59	55.50$\pm$0.58	–	5.22$\pm$0.20	38.20$\pm$0.78	10.00$\pm$0.32	15.60$\pm$0.64
9824	[C i]	–	1.71$\pm$0.16	–	–	–	2.41$\pm$0.76	0.98$\pm$0.15	0.82$\pm$0.09
9850	[C i]	–	1.67$\pm$0.16	–	–	2.49$\pm$0.17	3.32$\pm$0.76	2.58$\pm$0.15	1.74$\pm$0.09
10049	Pa$\delta$	–	–	–	–	–	–	3.85$\pm$0.44	–
10122	He ii	–	–	–	–	–	–	3.28$\pm$0.23	–
10830	He i	–	7.72$\pm$1.08	–	–	8.40$\pm$1.23	25.60$\pm$1.52	8.41$\pm$0.88	9.47$\pm$0.77
10938	Pa$\gamma$	–	3.98$\pm$1.08	–	–	2.57$\pm$0.57	8.07$\pm$1.20	3.02$\pm$0.47	4.08$\pm$0.55
11470	[P ii]	–	1.30$\pm$0.37	–	–	1.17$\pm$0.14	–	2.50$\pm$0.87	–
11886	[P ii]	–	1.80$\pm$0.37	3.83$\pm$1.08	–	1.46$\pm$0.14	–	4.98$\pm$0.87	–
12567	[Fe ii]	–	5.04$\pm$0.25	4.64$\pm$0.52	2.06$\pm$0.23	4.73$\pm$0.33	11.80$\pm$0.66	6.80$\pm$0.21	3.28$\pm$0.49
12820	Pa$\beta$	–	12.10$\pm$0.26	3.53$\pm$0.57	4.86$\pm$0.19	7.86$\pm$0.30	9.57$\pm$2.61	9.74$\pm$0.21	9.20$\pm$0.51
12950	[Fe ii]	–	1.18$\pm$0.42	–	–	–	–	–	–
13209	[Fe ii]	–	2.94$\pm$0.42	–	–	–	5.60$\pm$0.34	–	–
15342	[Fe ii]	–	–	–	–	–	–	–	–
16436	[Fe ii]	–	4.65$\pm$0.47	5.53$\pm$0.48	3.24$\pm$0.08	3.56$\pm$0.18	9.20$\pm$0.28	6.05$\pm$0.60	3.09$\pm$0.21
16773	[Fe ii]+Br11	–	–	–	–	–	–	–	–
17360	Br10	–	–	–	–	–	–	–	–
18750	Pa$\alpha$	4.70$\pm$0.18	60.40$\pm$0.46	82.50$\pm$3.49	32.60$\pm$0.29	34.30$\pm$0.17	31.60$\pm$3.36	36.70$\pm$0.67	23.30$\pm$0.32
19446	Br$\delta$	–	2.12$\pm$0.39	–	2.36$\pm$0.22	–	–	–	1.05$\pm$0.10
19570	H2	1.79$\pm$0.49	4.10$\pm$0.52	3.36$\pm$0.54	3.33$\pm$0.44	5.25$\pm$0.05	4.33$\pm$0.16	8.17$\pm$0.26	1.45$\pm$0.10
20332	H2	0.61$\pm$0.08	1.48$\pm$0.30	1.25$\pm$0.20	0.86$\pm$0.05	1.65$\pm$0.07	2.36$\pm$0.39	2.46$\pm$0.20	2.03$\pm$0.20
20580	H2	–	1.40$\pm$0.26	–	1.98$\pm$0.05	–	1.33$\pm$0.34	1.53$\pm$0.18	1.24$\pm$0.20
21218	H2	0.60$\pm$0.02	3.21$\pm$0.28	2.14$\pm$0.20	1.34$\pm$0.21	4.00$\pm$0.13	3.92$\pm$0.27	4.80$\pm$0.36	0.88$\pm$0.14
21654	Br$\gamma$	0.80$\pm$0.04	5.33$\pm$0.09	0.52$\pm$0.32	4.66$\pm$0.25	2.88$\pm$0.13	3.75$\pm$0.69	4.07$\pm$0.05	2.56$\pm$0.15
22230	H2	–	1.20$\pm$0.17	1.05$\pm$0.23	0.65$\pm$0.18	1.70$\pm$0.85	1.71$\pm$0.10	2.25$\pm$0.25	0.71$\pm$0.03
22470	H2	–	1.18$\pm$0.16	0.61$\pm$0.13	–	0.55$\pm$0.12	0.49$\pm$0.06	1.30$\pm$0.26	0.32$\pm$0.14

Table 5. Table 5: Line limits and continuum bandpasses.

Centre	Main Absorber	Index Name	Line Limits	Blue continuum	Red Continuum	Reference
(Å)			(Å)	(Å)	(Å)
4228.5	Ca i	Ca4227	4222.250 – 4234.750	4211.000 – 4219.750	4241.000 – 4251.000	Worthey et al. (1994)
4298.875	CH (G-Band)	G4300	4281.375 – 4316.375	4266.375 – 4282.625	4318.875 – 4335.125	Worthey et al. (1994)
4394.75	Fe i	Fe4383	4369.125 – 4420.375	4359.125 – 4370.375	4442.875 – 4455.375	Worthey et al. (1994)
4463.375	Ca i	Ca4455	4452.125 – 4474.625	4445.875 – 4454.625	4477.125 – 4492.125	Worthey et al. (1994)
4536.75	Fe i	Fe4531	4514.250 – 4559.250	4504.250 – 4514.250	4560.500 – 4579.250	Worthey et al. (1994)
4677.125	C2	Fe4668	4634.000 – 4720.250	4611.500 – 4630.250	4742.750 – 4756.500	Worthey et al. (1994)
5015.875	Fe i	Fe5015	4977.750 – 5054.000	4946.500 – 4977.750	5054.000 – 5065.250	Worthey et al. (1994)
5101.625	MgH	Mg1	5069.125 – 5134.125	4895.125 – 4957.625	5301.125 – 5366.125	Worthey et al. (1994)
5175.375	MgH	Mg2	5154.125 – 5196.625	4895.125 – 4957.625	5301.125 – 5366.125	Worthey et al. (1994)
5176.375	Mg b	Mgb	5160.125 – 5192.625	5142.625 – 5161.375	5191.375 – 5206.375	Worthey et al. (1994)
5265.65	Fe i	Fe5270	5245.650 – 5285.650	5233.150 – 5248.150	5285.650 – 5318.150	Worthey et al. (1994)
5332.125	Fe i	Fe5335	5312.125 – 5352.125	5304.625 – 5315.875	5353.375 – 5363.375	Worthey et al. (1994)
5401.25	Fe i	Fe5406	5387.500 – 5415.000	5376.250 – 5387.500	5415.000 – 5425.000	Worthey et al. (1994)
5708.5	Fe i	Fe5709	5696.625 – 5720.375	5672.875 – 5696.625	5722.875 – 5736.625	Worthey et al. (1994)
5786.625	Fe i	Fe5782	5776.625 – 5796.625	5765.375 – 5775.375	5797.875 – 5811.625	Worthey et al. (1994)
5893.125	Na i	NaD	5876.875 – 5909.375	5860.625 – 5875.625	5922.125 – 5948.125	Worthey et al. (1994)
5965.375	TiO	TiO1	5936.625 – 5994.125	5816.625 – 5849.125	6038.625 – 6103.625	Worthey et al. (1994)
6230.875	TiO	TiO2	6189.625 – 6272.125	6066.625 – 6141.625	6372.625 – 6415.125	Worthey et al. (1994)
8498.0	Ca ii	CaT1	8476.000 – 8520.000	8110.000 – 8165.000	8786.000 – 8844.000	Bica & Alloin (1987) ($†$)
8542.0	Ca ii	CaT2	8520.000 – 8564.000	8110.000 – 8165.000	8786.000 – 8844.000	Bica & Alloin (1987) ($†$)
8670.0	Ca ii	CaT3	8640.000 – 8700.000	8110.000 – 8165.000	8786.000 – 8844.000	Bica & Alloin (1987) ($†$)
9320.0	ZrO/TiO/CN	ZrO	9170.000 – 9470.000	8900.000 – 8960.000	9585.000 – 9615.000	New Definition ($\alpha$)
10560.0	VO	VO	10470.000 – 10650.000	10430.000 – 10465.000	10660.000 – 10700.000	New Definition ($\alpha$)
11000.0	CN	CN11	10910.000 – 11090.000	10705.000 – 10730.000	11310.000 – 11345.000	New Definition ($\beta$)
11390.0	Na i	NaI1.14	11350.000 – 11430.000	11310.000 – 11345.000	11450.000 – 11515.000	New Definition ($\beta$)
11605.0	Fe i	FeI1.16	11580.000 – 11630.000	11450.000 – 11515.000	11650.000 – 11690.000	Roeck (2015)
12430.0	Mg i	MgI1.24	12405.000 – 12455.000	12335.000 – 12365.000	12465.000 – 12490.000	Roeck (2015)
12944.0	Mn i	MnI1.29	12893.000 – 12995.000	12858.000 – 12878.000	13026.000 – 13068.000	New Definition
13132.5	Al i	AlI1.31	13095.000 – 13170.000	13000.000 – 13070.000	13175.000 – 13215.000	Roeck (2015)
14875.0	Mg i	MgI1.48	14850.000 – 14900.000	14750.000 – 14800.000	14910.000 – 14950.000	New Definition
15032.5	Mg i	MgI1.50	14995.000 – 15070.000	14910.000 – 14950.000	15150.000 – 15200.000	New Definition
15587.5	CO+Mg i	CO1.5a	15555.000 – 15620.000	15470.000 – 15500.000	15700.000 – 15730.000	New Definition ($ϵ$)
15780.0	CO+Mg i	CO1.5b	15750.000 – 15810.000	15700.000 – 15730.000	16095.000 – 16145.000	New Definition ($ϵ$)
15830.0	Fe i	FeI1.58	15810.000 – 15850.000	15700.000 – 15730.000	16090.000 – 16140.000	New Definition
15890.0	Si i + Mg i	SiI1.58	15850.000 – 15930.000	15700.000 – 15730.000	16090.000 – 16140.000	New Definition
15985.0	CO+Si i	CO1.5c	15950.000 – 16020.000	15700.000 – 15730.000	16090.000 – 16140.000	New Definition ($ϵ$)
16215.0	CO+Si i+Ca i	CO1.6a	16145.000 – 16285.000	16090.000 – 16140.000	16290.000 – 16340.000	New Definition ($ϵ$)
17064.0	CO + Fe i	CO1.6b	17035.000 – 17093.000	16970.000 – 17025.000	17140.000 – 17200.000	New Definition ($ϵ$)
17111.5	Mg i	MgI1.7	17093.000 – 17130.000	16970.000 – 17025.000	17140.000 – 17200.000	Roeck (2015)
22073.5	Na i	NaI2.20	22040.000 – 22107.000	21910.000 – 21966.000	22125.000 – 22160.000	Frogel et al. (2001)
22634.5	Ca i	CaI2.26	22577.000 – 22692.000	22530.000 – 22560.000	22700.000 – 22720.000	Frogel et al. (2001) ($\gamma$)
22820.0	Mg i	MgI2.28	22795.000 – 22845.000	22700.000 – 22720.000	22850.000 – 22865.000	New Definition ($\delta$)
23015.0	CO	CO2.2	22870.000 – 23160.000	22700.000 – 22790.000	23655.000 – 23680.000	New Definition ($ϵ$)
23290.0	CO	CO2.3a	23160.000 – 23420.000	22700.000 – 22790.000	23655.000 – 23680.000	New Definition ($ϵ$)
23535.0	CO	CO2.3b	23420.000 – 23650.000	22700.000 – 22790.000	23655.000 – 23680.000	New Definition ($ϵ$)

Table 6. Table 6: Absorption feature Equivalent Widths (in Å) .

Line	NGC23	NGC520	NGC660	NGC1055	NGC1134	NGC1204	NGC1222	NGC1266
Ca4227	0.36$\pm$0.06	–	0.53$\pm$0.61	–	1.14$\pm$0.16	–	–	–
G4300	1.56$\pm$0.19	–	5.18$\pm$1.62	–	2.01$\pm$0.56	–	–	–
Fe4383	1.67$\pm$0.22	–	7.23$\pm$0.68	–	3.84$\pm$1.05	–	–	7.53$\pm$2.22
Ca4455	0.49$\pm$0.1	3.22$\pm$0.47	2.01$\pm$0.3	–	0.28$\pm$0.22	–	–	–
Fe4531	2.26$\pm$0.16	–	–	–	1.2$\pm$0.46	–	–	–
$C_2$4668	3.88$\pm$0.2	–	–	–	4.94$\pm$0.55	–	–	–
Fe5015	–	–	–	–	–	–	–	–
Mg1	3.84$\pm$0.19	–	5.26$\pm$0.3	5.2$\pm$0.33	4.06$\pm$0.31	4.39$\pm$0.27	–	3.27$\pm$0.74
Mg2	4.98$\pm$0.13	–	6.46$\pm$0.21	6.73$\pm$0.2	5.36$\pm$0.18	5.24$\pm$0.18	–	4.14$\pm$0.37
Mgb	2.52$\pm$0.15	–	2.71$\pm$0.25	3.08$\pm$0.26	2.9$\pm$0.27	3.03$\pm$0.31	0.85$\pm$0.19	3.81$\pm$0.5
Fe5270	1.91$\pm$0.12	–	2.53$\pm$0.35	1.92$\pm$0.32	2.44$\pm$0.2	2.09$\pm$0.3	–	2.43$\pm$0.45
Fe5335	1.73$\pm$0.12	–	2.02$\pm$0.24	1.64$\pm$0.34	2.17$\pm$0.2	1.48$\pm$0.33	0.66$\pm$0.28	2.38$\pm$0.67
Fe5406	0.95$\pm$0.04	–	1.0$\pm$0.15	0.57$\pm$0.27	1.25$\pm$0.14	0.68$\pm$0.3	0.09$\pm$0.1	2.33$\pm$0.32
Fe5709	0.64$\pm$0.04	–	0.77$\pm$0.1	0.72$\pm$0.13	0.79$\pm$0.07	0.75$\pm$0.14	0.45$\pm$0.04	0.74$\pm$0.4
Fe5782	0.41$\pm$0.04	–	1.11$\pm$0.08	–	0.68$\pm$0.06	0.05$\pm$0.16	0.68$\pm$0.18	0.98$\pm$0.22
NaD	4.25$\pm$0.08	1.91$\pm$0.25	5.07$\pm$0.19	4.02$\pm$0.27	4.71$\pm$0.12	3.56$\pm$0.21	–	6.37$\pm$0.23
TiO1	0.57$\pm$0.09	–	–	–	–	0.51$\pm$0.24	–	–
TiO2	3.82$\pm$0.11	–	5.79$\pm$0.2	7.96$\pm$0.28	6.13$\pm$0.2	6.07$\pm$0.27	3.87$\pm$0.19	8.34$\pm$0.51
CaT1	4.13$\pm$0.11	–	–	–	3.62$\pm$0.20	1.16$\pm$0.36	3.93$\pm$0.13	5.74$\pm$1.14
CaT2	5.48$\pm$0.09	–	–	–	7.51$\pm$0.17	3.11$\pm$0.31	5.46$\pm$0.13	6.16$\pm$1.13
CaT3	3.22$\pm$0.16	–	–	–	3.07$\pm$0.50	–	2.97$\pm$0.17	–
ZrO	16.76$\pm$0.27	–	–	–	13.09$\pm$1.29	15.60$\pm$1.78	14.06$\pm$0.56	6.70$\pm$2.53
VO	0.05$\pm$0.31	–	1.41$\pm$0.63	–	–	–	–	–
CN11	12.32$\pm$0.14	–	3.76$\pm$0.28	–	11.15$\pm$0.17	6.48$\pm$0.31	6.37$\pm$0.27	12.41$\pm$0.92
NaI1.14	1.74$\pm$0.08	–	2.43$\pm$0.18	–	1.48$\pm$0.11	1.34$\pm$0.18	2.01$\pm$0.06	3.46$\pm$0.44
FeI1.16	0.71$\pm$0.05	–	0.44$\pm$0.07	–	0.21$\pm$0.06	0.71$\pm$0.07	0.65$\pm$0.06	–
MgI1.24	0.99$\pm$0.06	–	0.74$\pm$0.05	–	1.68$\pm$0.10	1.15$\pm$0.07	–	0.57$\pm$0.15
MnI1.29	0.03$\pm$0.15	–	0.28$\pm$0.14	–	0.06$\pm$0.25	0.80$\pm$0.35	0.55$\pm$0.10	3.32$\pm$0.14
AlI1.31	1.54$\pm$0.07	–	1.93$\pm$0.57	–	2.17$\pm$0.10	2.16$\pm$0.34	1.90$\pm$0.07	3.04$\pm$0.15
MgI1.48	1.80$\pm$0.03	2.94$\pm$0.17	1.96$\pm$0.03	1.14$\pm$0.25	1.67$\pm$0.04	1.16$\pm$0.06	1.07$\pm$0.07	1.15$\pm$0.11
MgI1.50	3.77$\pm$0.07	3.28$\pm$0.30	2.43$\pm$0.10	–	4.35$\pm$0.09	3.46$\pm$0.08	2.44$\pm$0.08	2.81$\pm$0.13
CO1.5a	3.52$\pm$0.09	6.51$\pm$0.23	4.33$\pm$0.10	5.12$\pm$0.36	2.61$\pm$0.23	3.18$\pm$0.13	2.66$\pm$0.05	5.24$\pm$0.15
CO1.5b	4.26$\pm$0.11	6.76$\pm$0.23	4.94$\pm$0.10	4.71$\pm$0.19	4.44$\pm$0.19	4.35$\pm$0.08	2.50$\pm$0.05	5.28$\pm$0.07
FeI1.58	1.50$\pm$0.06	3.64$\pm$0.13	2.11$\pm$0.06	0.45$\pm$0.10	0.89$\pm$0.11	1.85$\pm$0.06	1.01$\pm$0.04	1.66$\pm$0.05
SiI1.58	3.66$\pm$0.10	3.65$\pm$0.24	4.03$\pm$0.13	4.40$\pm$0.20	3.77$\pm$0.18	4.25$\pm$0.11	3.00$\pm$0.08	4.63$\pm$0.10
CO1.5c	3.50$\pm$0.06	3.39$\pm$0.20	3.83$\pm$0.11	3.16$\pm$0.17	2.72$\pm$0.10	4.07$\pm$0.11	2.22$\pm$0.12	4.49$\pm$0.12
CO1.6a	4.97$\pm$0.11	7.60$\pm$0.39	6.70$\pm$0.15	7.21$\pm$0.44	5.43$\pm$0.16	7.73$\pm$0.25	4.29$\pm$0.29	4.73$\pm$0.30
CO1.6b	2.32$\pm$0.05	3.21$\pm$0.19	0.79$\pm$0.07	1.11$\pm$0.62	1.57$\pm$0.08	1.34$\pm$0.13	1.75$\pm$0.07	2.92$\pm$0.22
MgI1.7	1.68$\pm$0.05	0.76$\pm$0.21	1.37$\pm$0.04	0.94$\pm$0.45	0.91$\pm$0.04	1.35$\pm$0.08	1.16$\pm$0.04	0.23$\pm$0.21
NaI2.20	3.45$\pm$0.08	2.51$\pm$0.10	2.88$\pm$0.07	4.26$\pm$0.26	2.82$\pm$0.08	3.24$\pm$0.04	1.29$\pm$0.08	3.31$\pm$0.12
CaI2.26	3.07$\pm$0.07	4.07$\pm$0.12	2.40$\pm$0.18	3.87$\pm$0.39	1.78$\pm$0.07	2.21$\pm$0.14	2.25$\pm$0.10	2.15$\pm$0.08
MgI2.28	1.14$\pm$0.02	1.20$\pm$0.09	0.60$\pm$0.04	5.57$\pm$0.10	0.62$\pm$0.04	0.37$\pm$0.04	0.31$\pm$0.06	0.03$\pm$0.01
CO2.2	20.26$\pm$0.56	23.16$\pm$0.36	12.23$\pm$0.39	23.11$\pm$1.22	19.33$\pm$0.41	21.80$\pm$0.58	18.03$\pm$0.61	24.34$\pm$0.61
CO2.3a	19.07$\pm$0.38	26.71$\pm$0.35	12.45$\pm$0.66	22.69$\pm$0.80	21.93$\pm$0.26	22.02$\pm$0.43	18.71$\pm$0.64	24.94$\pm$0.40
CO2.3b	21.53$\pm$0.37	27.00$\pm$0.42	12.50$\pm$0.78	18.28$\pm$0.56	24.03$\pm$0.13	22.57$\pm$0.44	13.99$\pm$0.85	24.52$\pm$0.26

Table 7. Table 7: Absorption feature Equivalent Widths (in Å).

Line	UGC2982	NGC1797	NGC6814	NGC6835	UGC12150	NGC7465	NGC7591	NGC7678
Ca4227	–	0.91$\pm$0.12	–	–	–	0.47$\pm$0.1	–	–
G4300	–	1.65$\pm$0.44	–	–	–	2.84$\pm$0.52	–	0.78$\pm$0.34
Fe4383	–	–	–	–	–	2.77$\pm$0.27	–	2.24$\pm$0.4
Ca4455	–	–	–	–	–	0.53$\pm$0.16	–	1.11$\pm$0.12
Fe4531	–	1.93$\pm$0.37	–	–	–	2.67$\pm$0.16	–	1.48$\pm$0.32
$C_2$4668	–	–	–	–	–	–	–	1.1$\pm$0.42
Fe5015	–	–	–	–	–	–	–	–
Mg1	–	2.35$\pm$0.21	–	–	–	–	–	3.61$\pm$0.33
Mg2	–	3.77$\pm$0.13	–	–	4.28$\pm$0.26	4.07$\pm$0.24	5.14$\pm$0.36	4.23$\pm$0.16
Mgb	–	2.12$\pm$0.24	–	–	2.04$\pm$0.27	2.56$\pm$0.23	3.09$\pm$0.28	1.96$\pm$0.23
Fe5270	–	1.32$\pm$0.19	–	–	1.34$\pm$0.28	1.78$\pm$0.17	1.82$\pm$0.29	1.31$\pm$0.16
Fe5335	–	1.22$\pm$0.13	–	–	1.11$\pm$0.37	1.85$\pm$0.15	2.32$\pm$0.3	1.5$\pm$0.29
Fe5406	–	0.87$\pm$0.11	–	–	0.44$\pm$0.2	0.78$\pm$0.15	1.31$\pm$0.21	1.04$\pm$0.12
Fe5709	–	0.35$\pm$0.06	–	–	0.64$\pm$0.1	0.7$\pm$0.06	1.14$\pm$0.17	0.78$\pm$0.11
Fe5782	–	0.31$\pm$0.06	–	–	–	0.57$\pm$0.03	0.61$\pm$0.14	0.86$\pm$0.09
NaD	4.78$\pm$0.37	6.27$\pm$0.19	–	–	7.47$\pm$0.52	1.67$\pm$0.27	3.86$\pm$0.19	4.17$\pm$0.25
TiO1	–	–	–	–	–	1.4$\pm$0.16	–	–
TiO2	–	2.08$\pm$0.23	–	–	4.88$\pm$0.52	4.63$\pm$0.15	3.4$\pm$0.25	4.6$\pm$0.21
CaT1	–	1.13$\pm$0.39	–	–	2.22$\pm$0.42	5.17$\pm$0.17	0.80$\pm$0.33	2.73$\pm$0.26
CaT2	–	3.16$\pm$0.34	0.36$\pm$0.19	–	1.55$\pm$0.43	6.28$\pm$0.14	5.37$\pm$0.26	4.29$\pm$0.24
CaT3	–	–	3.35$\pm$0.06	–	–	5.93$\pm$0.33	3.73$\pm$0.39	5.90$\pm$0.26
ZrO	–	15.11$\pm$2.04	1.20$\pm$1.74	–	12.38$\pm$0.98	9.16$\pm$0.75	14.76$\pm$1.61	6.18$\pm$1.10
VO	7.10$\pm$1.35	–	3.89$\pm$0.43	–	1.13$\pm$0.43	–	–	–
CN11	20.92$\pm$0.74	6.49$\pm$0.33	–	–	7.95$\pm$0.42	6.91$\pm$0.37	11.93$\pm$0.34	5.08$\pm$0.72
NaI1.14	–	1.38$\pm$0.19	3.99$\pm$0.13	7.92$\pm$0.83	1.75$\pm$0.15	1.66$\pm$0.18	1.17$\pm$0.22	0.78$\pm$0.11
FeI1.16	0.17$\pm$0.20	0.70$\pm$0.07	0.53$\pm$0.06	2.78$\pm$0.11	0.68$\pm$0.08	0.76$\pm$0.06	1.25$\pm$0.06	–
MgI1.24	1.40$\pm$0.06	1.16$\pm$0.07	0.76$\pm$0.05	1.49$\pm$0.06	0.71$\pm$0.15	0.63$\pm$0.04	0.67$\pm$0.02	0.41$\pm$0.06
MnI1.29	3.11$\pm$0.62	0.77$\pm$0.33	10.00$\pm$0.43	1.53$\pm$0.11	1.74$\pm$0.14	1.12$\pm$0.10	–	2.18$\pm$0.20
AlI1.31	0.46$\pm$0.26	2.15$\pm$0.35	1.33$\pm$0.11	–	3.16$\pm$0.34	2.29$\pm$0.15	2.89$\pm$0.16	2.84$\pm$0.17
MgI1.48	0.57$\pm$0.10	1.16$\pm$0.07	0.94$\pm$0.03	1.83$\pm$0.14	0.80$\pm$0.10	1.83$\pm$0.09	1.38$\pm$0.05	1.77$\pm$0.08
MgI1.50	2.41$\pm$0.10	3.46$\pm$0.08	2.07$\pm$0.07	1.82$\pm$0.25	2.88$\pm$0.18	3.12$\pm$0.17	3.11$\pm$0.14	3.53$\pm$0.14
CO1.5a	2.83$\pm$0.05	3.21$\pm$0.10	2.53$\pm$0.07	3.06$\pm$0.07	4.03$\pm$0.11	3.89$\pm$0.07	4.39$\pm$0.13	2.98$\pm$0.09
CO1.5b	3.46$\pm$0.04	4.34$\pm$0.07	3.10$\pm$0.05	4.63$\pm$0.19	5.34$\pm$0.13	3.48$\pm$0.09	4.55$\pm$0.10	3.52$\pm$0.06
FeI1.58	0.73$\pm$0.03	1.85$\pm$0.05	0.77$\pm$0.03	2.46$\pm$0.11	1.80$\pm$0.09	1.35$\pm$0.06	1.81$\pm$0.07	1.18$\pm$0.05
SiI1.58	2.81$\pm$0.07	4.25$\pm$0.12	2.57$\pm$0.07	4.45$\pm$0.19	4.58$\pm$0.17	3.43$\pm$0.11	3.71$\pm$0.15	2.01$\pm$0.14
CO1.5c	2.99$\pm$0.07	4.09$\pm$0.10	2.51$\pm$0.06	3.46$\pm$0.18	3.39$\pm$0.13	2.93$\pm$0.13	4.03$\pm$0.19	2.70$\pm$0.16
CO1.6a	6.02$\pm$0.19	7.73$\pm$0.27	4.26$\pm$0.20	9.42$\pm$0.45	7.24$\pm$0.18	6.17$\pm$0.20	7.73$\pm$0.26	6.50$\pm$0.28
CO1.6b	2.64$\pm$0.13	1.36$\pm$0.12	0.92$\pm$0.05	2.07$\pm$0.28	–	1.13$\pm$0.09	1.66$\pm$0.10	–
MgI1.7	0.85$\pm$0.11	1.35$\pm$0.09	1.11$\pm$0.03	2.21$\pm$0.29	1.34$\pm$0.14	1.82$\pm$0.05	1.76$\pm$0.07	1.73$\pm$0.04
NaI2.20	4.15$\pm$0.15	3.24$\pm$0.05	1.54$\pm$0.03	3.41$\pm$0.04	4.10$\pm$0.07	2.70$\pm$0.04	3.63$\pm$0.08	2.88$\pm$0.07
CaI2.26	6.74$\pm$0.13	2.21$\pm$0.16	1.04$\pm$0.04	2.43$\pm$0.04	4.37$\pm$0.25	1.87$\pm$0.09	4.35$\pm$0.11	2.49$\pm$0.13
MgI2.28	0.71$\pm$0.18	0.36$\pm$0.04	0.05$\pm$0.04	1.05$\pm$0.04	–	0.33$\pm$0.05	1.58$\pm$0.05	0.86$\pm$0.06
CO2.2	17.96$\pm$0.81	21.54$\pm$0.56	6.24$\pm$0.12	22.26$\pm$0.54	22.08$\pm$1.02	16.14$\pm$0.67	23.89$\pm$0.47	17.53$\pm$0.83
CO2.3a	22.66$\pm$1.19	22.13$\pm$0.39	2.96$\pm$0.16	23.54$\pm$0.41	22.84$\pm$0.95	14.90$\pm$0.82	26.34$\pm$0.40	24.72$\pm$0.77
CO2.3b	27.53$\pm$1.39	22.58$\pm$0.44	4.36$\pm$0.20	24.65$\pm$0.41	20.18$\pm$1.07	14.97$\pm$1.47	27.79$\pm$0.50	20.27$\pm$0.89

Table 8. Table 8: Absorption feature Equivalent Widths (in Å) from the sample of Baldwin et al. ( 2017 ) .

Line	IC0719	NGC3032	NGC3098	NGC3156	NGC3182	NGC3301
Ca4227	1.68$\pm$0.01	–	–	0.56$\pm$0.00	1.41$\pm$0.00	–
G4300	6.23$\pm$0.04	–	–	1.94$\pm$0.00	5.57$\pm$0.01	–
Fe4383	5.99$\pm$0.05	–	–	1.89$\pm$0.00	5.33$\pm$0.01	–
Ca4455	2.07$\pm$0.02	–	–	1.08$\pm$0.00	1.68$\pm$0.00	–
Fe4531	3.67$\pm$0.04	–	–	2.99$\pm$0.00	3.37$\pm$0.01	–
$C_2$4668	7.04$\pm$0.07	–	–	2.79$\pm$0.01	7.86$\pm$0.01	–
Fe5015	5.68$\pm$0.06	4.44$\pm$0.14	4.64$\pm$0.14	4.07$\pm$0.01	3.57$\pm$0.01	4.95$\pm$0.14
Mg1	5.40$\pm$0.05	–	–	1.88$\pm$0.01	6.46$\pm$0.01	–
Mg2	8.06$\pm$0.03	–	–	4.10$\pm$0.00	8.82$\pm$0.01	–
Mgb	3.83$\pm$0.03	1.70$\pm$0.12	3.17$\pm$0.12	1.58$\pm$0.00	4.18$\pm$0.01	2.93$\pm$0.12
Fe5270	3.02$\pm$0.03	2.06$\pm$0.11	2.56$\pm$0.11	1.93$\pm$0.01	2.97$\pm$0.01	2.56$\pm$0.11
Fe5335	2.93$\pm$0.02	–	–	1.98$\pm$0.00	3.04$\pm$0.01	–
Fe5406	1.51$\pm$0.02	–	–	1.06$\pm$0.00	1.87$\pm$0.00	–
Fe5709	1.21$\pm$0.01	–	–	0.68$\pm$0.00	1.09$\pm$0.00	–
Fe5782	0.97$\pm$0.01	–	–	0.51$\pm$0.00	0.91$\pm$0.00	–
NaD	3.13$\pm$0.03	–	–	1.19$\pm$0.01	3.63$\pm$0.01	–
TiO1	2.20$\pm$0.04	–	–	0.60$\pm$0.01	1.66$\pm$0.01	–
TiO2	5.17$\pm$0.05	–	–	1.87$\pm$0.02	5.26$\pm$0.02	–
CaT1	–	–	–	–	–	–
CaT2	–	–	–	–	–	–
CaT3	–	–	–	–	–	–
ZrO	8.55$\pm$2.34	15.87$\pm$1.08	8.92$\pm$1.41	22.28$\pm$1.91	14.75$\pm$1.68	15.61$\pm$0.92
VO	2.63$\pm$0.55	–	0.32$\pm$0.51	1.87$\pm$0.59	0.28$\pm$0.38	1.45$\pm$0.39
CN11	5.97$\pm$0.79	10.58$\pm$0.37	11.35$\pm$0.39	7.31$\pm$0.64	7.93$\pm$0.58	8.47$\pm$0.36
NaI1.14	1.82$\pm$0.21	3.14$\pm$0.17	1.48$\pm$0.14	1.23$\pm$0.26	2.37$\pm$0.17	1.42$\pm$0.14
FeI1.16	1.23$\pm$0.26	–	0.61$\pm$0.10	0.45$\pm$0.18	0.39$\pm$0.22	0.03$\pm$0.05
MgI1.24	1.00$\pm$0.12	1.04$\pm$0.09	0.36$\pm$0.08	0.81$\pm$0.15	1.43$\pm$0.11	0.81$\pm$0.07
MnI1.29	0.85$\pm$0.19	1.06$\pm$0.09	1.20$\pm$0.13	1.19$\pm$0.20	1.20$\pm$0.20	1.26$\pm$0.18
AlI1.31	1.27$\pm$0.16	1.87$\pm$0.10	1.32$\pm$0.18	0.74$\pm$0.15	2.27$\pm$0.11	1.27$\pm$0.11
MgI1.48	2.05$\pm$0.13	1.62$\pm$0.11	1.71$\pm$0.08	1.12$\pm$0.13	0.85$\pm$0.27	1.55$\pm$0.09
MgI1.50	4.29$\pm$0.23	3.58$\pm$0.17	3.60$\pm$0.21	3.23$\pm$0.22	1.42$\pm$0.31	3.37$\pm$0.18
CO1.5a	3.72$\pm$0.14	4.74$\pm$0.16	4.05$\pm$0.28	4.40$\pm$0.19	2.75$\pm$0.27	4.38$\pm$0.18
CO1.5b	4.13$\pm$0.20	4.09$\pm$0.14	4.96$\pm$0.20	3.61$\pm$0.16	2.61$\pm$0.44	5.56$\pm$0.10
FeI1.58	1.37$\pm$0.13	0.85$\pm$0.11	1.55$\pm$0.11	1.78$\pm$0.11	0.62$\pm$0.24	2.14$\pm$0.07
SiI1.58	2.91$\pm$0.22	3.44$\pm$0.24	3.24$\pm$0.23	4.14$\pm$0.23	2.89$\pm$0.49	4.85$\pm$0.14
CO1.5c	3.97$\pm$0.17	3.56$\pm$0.24	3.20$\pm$0.16	2.83$\pm$0.22	3.35$\pm$0.29	3.98$\pm$0.16
CO1.6a	5.33$\pm$0.35	4.73$\pm$0.45	6.67$\pm$0.30	6.48$\pm$0.41	5.81$\pm$0.48	6.30$\pm$0.37
CO1.6b	1.27$\pm$0.16	0.89$\pm$0.14	1.51$\pm$0.13	1.28$\pm$0.16	1.37$\pm$0.21	1.31$\pm$0.15
MgI1.7	1.99$\pm$0.08	2.30$\pm$0.08	2.12$\pm$0.09	1.85$\pm$0.10	2.41$\pm$0.09	2.14$\pm$0.10
NaI2.20	3.41$\pm$0.08	5.15$\pm$0.14	2.84$\pm$0.13	2.52$\pm$0.20	3.74$\pm$0.08	2.77$\pm$0.19
CaI2.26	1.85$\pm$0.17	3.66$\pm$0.26	2.49$\pm$0.34	2.55$\pm$0.29	2.57$\pm$0.22	3.09$\pm$0.14
MgI2.28	0.76$\pm$0.11	0.94$\pm$0.12	1.05$\pm$0.09	–	1.22$\pm$0.16	0.94$\pm$0.04
CO2.2	15.76$\pm$0.83	20.43$\pm$0.68	14.80$\pm$0.70	19.55$\pm$0.78	17.57$\pm$0.77	19.47$\pm$0.58
CO2.3a	17.74$\pm$1.01	18.16$\pm$0.65	14.12$\pm$0.86	19.90$\pm$1.21	16.44$\pm$0.78	19.23$\pm$0.70
CO2.3b	20.23$\pm$1.27	17.06$\pm$0.78	17.40$\pm$1.13	23.39$\pm$1.85	20.04$\pm$1.15	21.83$\pm$0.98

Table 9. Table 9: Absorption feature Equivalent Widths (in Å) from the sample of Baldwin et al. ( 2017 ) .

Line	NGC3489	NGC4379	NGC4578	NGC4608	NGC4710	NGC5475
Ca4227	–	–	–	–	–	1.32$\pm$0.00
G4300	–	–	–	–	–	5.76$\pm$0.01
Fe4383	–	–	–	–	–	5.43$\pm$0.01
Ca4455	–	–	–	–	–	1.87$\pm$0.00
Fe4531	–	–	–	–	–	3.79$\pm$0.01
$C_2$4668	–	–	–	–	–	7.12$\pm$0.01
Fe5015	4.84$\pm$0.14	5.41$\pm$0.14	5.24$\pm$0.14	4.64$\pm$0.14	4.46$\pm$0.15	5.09$\pm$0.01
Mg1	–	–	–	–	–	5.33$\pm$0.01
Mg2	–	–	–	–	–	8.28$\pm$0.00
Mgb	2.80$\pm$0.12	3.66$\pm$0.12	3.88$\pm$0.12	3.78$\pm$0.12	3.05$\pm$0.1 2	4.22$\pm$0.00
Fe5270	2.58$\pm$0.11	2.85$\pm$0.12	2.90$\pm$0.11	2.59$\pm$0.11	2.44$\pm$0.15	3.10$\pm$0.00
Fe5335	–	–	–	–	–	3.24$\pm$0.00
Fe5406	–	–	–	–	–	1.96$\pm$0.00
Fe5709	–	–	–	–	–	1.09$\pm$0.00
Fe5782	–	–	–	–	–	1.07$\pm$0.00
NaD	–	–	–	–	–	3.58$\pm$0.00
TiO1	–	–	–	–	–	1.26$\pm$0.01
TiO2	–	–	–	–	–	4.73$\pm$0.01
CaT1	–	–	–	–	–	–
CaT2	–	–	–	–	–	–
CaT3	–	–	–	–	–	–
ZrO	12.10$\pm$0.68	25.45$\pm$0.64	22.91$\pm$0.98	19.96$\pm$1.62	6.11$\pm$1.51	6.40$\pm$1.26
VO	1.34$\pm$0.44	1.52$\pm$0.58	1.87$\pm$0.36	–	3.32$\pm$1.07	2.22$\pm$0.57
CN11	7.14$\pm$0.40	4.53$\pm$0.49	5.49$\pm$0.37	7.03$\pm$0.43	9.21$\pm$1.01	10.28$\pm$0.47
NaI1.14	1.35$\pm$0.20	1.47$\pm$0.19	1.63$\pm$0.16	1.69$\pm$0.21	1.22$\pm$0.25	1.61$\pm$0.21
FeI1.16	0.48$\pm$0.07	0.38$\pm$0.10	0.49$\pm$0.10	0.44$\pm$0.09	0.73$\pm$0.12	0.33$\pm$0.11
MgI1.24	0.60$\pm$0.04	0.88$\pm$0.06	0.99$\pm$0.07	0.94$\pm$0.07	0.71$\pm$0.15	0.80$\pm$0.08
MnI1.29	0.91$\pm$0.11	1.06$\pm$0.35	0.98$\pm$0.26	1.10$\pm$0.21	0.70$\pm$0.28	1.71$\pm$0.25
AlI1.31	1.19$\pm$0.10	1.35$\pm$0.18	1.65$\pm$0.12	1.65$\pm$0.09	1.38$\pm$0.15	2.21$\pm$0.10
MgI1.48	1.49$\pm$0.09	1.62$\pm$0.11	2.29$\pm$0.22	1.33$\pm$0.11	1.70$\pm$0.13	1.72$\pm$0.09
MgI1.50	3.13$\pm$0.13	4.07$\pm$0.18	5.10$\pm$0.35	3.86$\pm$0.15	4.02$\pm$0.19	4.54$\pm$0.21
CO1.5a	4.19$\pm$0.14	3.36$\pm$0.17	5.15$\pm$0.26	3.43$\pm$0.16	3.41$\pm$0.19	4.02$\pm$0.17
CO1.5b	4.46$\pm$0.13	3.10$\pm$0.19	5.95$\pm$0.47	3.86$\pm$0.16	3.44$\pm$0.21	4.55$\pm$0.15
FeI1.58	2.28$\pm$0.10	1.40$\pm$0.11	2.66$\pm$0.23	1.97$\pm$0.10	1.98$\pm$0.15	1.84$\pm$0.10
SiI1.58	3.75$\pm$0.18	2.97$\pm$0.18	4.43$\pm$0.42	3.98$\pm$0.21	3.70$\pm$0.22	4.03$\pm$0.20
CO1.5c	3.72$\pm$0.17	2.68$\pm$0.11	3.71$\pm$0.36	3.37$\pm$0.25	3.59$\pm$0.20	4.26$\pm$0.20
CO1.6a	6.74$\pm$0.29	5.51$\pm$0.29	6.64$\pm$0.54	6.78$\pm$0.37	5.59$\pm$0.34	6.07$\pm$0.37
CO1.6b	0.85$\pm$0.15	0.95$\pm$0.09	1.47$\pm$0.22	2.12$\pm$0.17	0.92$\pm$0.12	1.48$\pm$0.17
MgI1.7	1.94$\pm$0.06	2.25$\pm$0.06	2.27$\pm$0.10	2.33$\pm$0.08	2.13$\pm$0.06	2.32$\pm$0.11
NaI2.20	3.22$\pm$0.04	3.91$\pm$0.22	4.59$\pm$0.12	3.11$\pm$0.11	3.56$\pm$0.06	3.96$\pm$0.08
CaI2.26	2.38$\pm$0.09	0.88$\pm$0.42	2.10$\pm$0.15	2.09$\pm$0.30	2.41$\pm$0.28	3.40$\pm$0.30
MgI2.28	0.56$\pm$0.02	0.48$\pm$0.18	0.88$\pm$0.10	1.04$\pm$0.10	0.86$\pm$0.14	1.57$\pm$0.20
CO2.2	18.62$\pm$0.44	14.95$\pm$0.78	18.82$\pm$0.49	17.66$\pm$0.62	18.84$\pm$0.61	14.87$\pm$0.50
CO2.3a	17.39$\pm$0.53	14.74$\pm$0.81	19.96$\pm$0.67	16.66$\pm$0.75	18.11$\pm$0.63	16.69$\pm$0.68
CO2.3b	18.45$\pm$0.66	16.78$\pm$1.03	23.74$\pm$0.91	19.23$\pm$1.08	20.31$\pm$0.87	19.88$\pm$0.89

Table 10. Table 10: Absorption feature Equivalent Widths (in Å) from the sample of Dahmer-Hahn et al. ( 2018 ) .

Line	NGC4636	NGC5905	NGC5966	NGC6081	NGC6146	NGC6338	UGC08234
Ca4227	–	–	1.23$\pm$0.10	1.21$\pm$0.07	1.00$\pm$0.03	1.16$\pm$0.06	0.76$\pm$0.05
G4300	–	–	5.54$\pm$0.28	5.59$\pm$0.30	5.30$\pm$0.16	5.82$\pm$0.23	3.01$\pm$0.21
Fe4383	–	–	4.46$\pm$0.35	4.56$\pm$0.29	4.47$\pm$0.24	4.33$\pm$0.29	2.63$\pm$0.25
Ca4455	–	–	1.36$\pm$0.16	1.14$\pm$0.15	1.04$\pm$0.13	1.30$\pm$0.17	0.92$\pm$0.15
Fe4531	–	–	3.32$\pm$0.14	3.04$\pm$0.21	3.15$\pm$0.14	3.24$\pm$0.12	2.73$\pm$0.08
$C_2$4668	–	–	6.42$\pm$0.13	6.63$\pm$0.16	6.82$\pm$0.18	7.78$\pm$0.23	4.56$\pm$0.13
Fe5015	–	–	4.01$\pm$0.29	3.52$\pm$0.37	3.75$\pm$0.25	3.98$\pm$0.24	4.14$\pm$0.22
Mg1	–	–	6.28$\pm$0.27	6.84$\pm$0.28	6.58$\pm$0.22	8.06$\pm$0.22	3.26$\pm$0.20
Mg2	–	–	8.70$\pm$0.15	8.90$\pm$0.15	9.08$\pm$0.12	10.34$\pm$0.13	5.96$\pm$0.13
Mgb	–	–	3.95$\pm$0.15	4.09$\pm$0.15	3.89$\pm$0.13	4.67$\pm$0.12	2.60$\pm$0.09
Fe5270	–	–	2.51$\pm$0.15	2.53$\pm$0.14	2.47$\pm$0.10	2.40$\pm$0.12	2.30$\pm$0.15
Fe5335	–	–	2.19$\pm$0.13	2.19$\pm$0.13	1.97$\pm$0.08	2.15$\pm$0.11	1.92$\pm$0.10
Fe5406	–	–	1.62$\pm$0.06	1.36$\pm$0.05	1.34$\pm$0.08	1.46$\pm$0.05	1.27$\pm$0.04
Fe5709	–	–	0.85$\pm$0.04	0.71$\pm$0.05	0.73$\pm$0.04	0.26$\pm$0.07	0.57$\pm$0.12
Fe5782	–	–	0.50$\pm$0.03	0.82$\pm$0.03	0.67$\pm$0.03	0.73$\pm$0.03	0.60$\pm$0.03
NaD	–	–	3.45$\pm$0.06	4.02$\pm$0.04	4.00$\pm$0.03	4.82$\pm$0.02	3.18$\pm$0.04
TiO1	–	–	1.78$\pm$0.06	1.82$\pm$0.08	1.79$\pm$0.11	2.16$\pm$0.07	1.51$\pm$0.08
TiO2	–	–	5.80$\pm$0.17	5.96$\pm$0.09	5.52$\pm$0.15	5.45$\pm$0.15	3.56$\pm$0.14
CaT1	–	–	–	–	–	–	–
CaT2	–	–	–	–	–	–	–
CaT3	–	–	–	–	–	–	–
ZrO	–	–	–	–	–	–	–
VO	2.61$\pm$0.20	1.86$\pm$0.46	2.98$\pm$0.34	1.37$\pm$0.35	2.11$\pm$0.28	0.97$\pm$0.20	2.60$\pm$0.25
CN11	10.82$\pm$0.68	0.68$\pm$0.57	0.31$\pm$0.35	0.34$\pm$0.17	–	5.39$\pm$0.33	2.71$\pm$0.67
NaI1.14	1.57$\pm$0.47	0.94$\pm$0.24	–	1.18$\pm$0.09	2.21$\pm$0.17	1.15$\pm$0.09	1.70$\pm$0.26
FeI1.16	–	1.02$\pm$0.07	1.43$\pm$0.09	0.40$\pm$0.08	0.56$\pm$0.03	0.29$\pm$0.06	–
MgI1.24	0.56$\pm$0.06	1.18$\pm$0.13	1.12$\pm$0.06	0.93$\pm$0.17	0.47$\pm$0.05	2.06$\pm$0.14	0.76$\pm$0.26
MnI1.29	1.15$\pm$0.14	–	3.89$\pm$0.32	0.99$\pm$0.32	1.40$\pm$0.11	2.98$\pm$0.46	1.74$\pm$0.24
AlI1.31	3.54$\pm$0.19	1.61$\pm$0.06	2.58$\pm$0.16	2.98$\pm$0.14	0.02$\pm$0.04	0.56$\pm$0.07	–
MgI1.48	0.92$\pm$0.14	2.60$\pm$0.07	1.68$\pm$0.08	1.38$\pm$0.05	1.33$\pm$0.06	1.69$\pm$0.10	1.31$\pm$0.10
MgI1.50	3.88$\pm$0.12	3.92$\pm$0.12	3.50$\pm$0.17	4.03$\pm$0.05	3.03$\pm$0.09	3.24$\pm$0.17	3.23$\pm$0.17
CO1.5a	4.59$\pm$0.06	3.63$\pm$0.14	2.93$\pm$0.10	1.96$\pm$0.08	3.89$\pm$0.02	5.39$\pm$0.15	3.07$\pm$0.13
CO1.5b	3.89$\pm$0.08	4.73$\pm$0.18	4.91$\pm$0.09	3.61$\pm$0.11	3.95$\pm$0.03	4.52$\pm$0.26	4.56$\pm$0.09
FeI1.58	1.10$\pm$0.05	1.99$\pm$0.09	1.83$\pm$0.06	1.54$\pm$0.09	1.80$\pm$0.03	1.70$\pm$0.15	1.77$\pm$0.06
SiI1.58	2.79$\pm$0.07	3.49$\pm$0.18	4.90$\pm$0.11	3.20$\pm$0.15	3.89$\pm$0.06	4.20$\pm$0.25	4.17$\pm$0.12
CO1.5c	3.24$\pm$0.04	3.27$\pm$0.15	3.02$\pm$0.13	3.53$\pm$0.15	3.03$\pm$0.07	4.32$\pm$0.12	4.26$\pm$0.11
CO1.6a	7.43$\pm$0.17	5.99$\pm$0.30	8.47$\pm$0.24	6.18$\pm$0.26	5.60$\pm$0.14	6.30$\pm$0.12	5.85$\pm$0.20
CO1.6b	1.22$\pm$0.04	1.21$\pm$0.10	0.96$\pm$0.07	1.28$\pm$0.07	1.18$\pm$0.08	1.37$\pm$0.10	0.34$\pm$0.10
MgI1.7	2.12$\pm$0.02	1.81$\pm$0.06	2.14$\pm$0.03	1.81$\pm$0.05	2.18$\pm$0.05	1.76$\pm$0.06	1.57$\pm$0.08
NaI2.20	3.73$\pm$0.06	3.96$\pm$0.07	4.36$\pm$0.11	2.99$\pm$0.16	4.00$\pm$0.18	4.07$\pm$0.14	1.79$\pm$0.12
CaI2.26	3.83$\pm$0.04	2.16$\pm$0.46	0.81$\pm$0.39	2.07$\pm$0.13	2.51$\pm$0.29	5.05$\pm$0.07	5.31$\pm$0.41
MgI2.28	0.92$\pm$0.08	0.83$\pm$0.12	0.86$\pm$0.09	0.91$\pm$0.24	0.42$\pm$0.15	1.46$\pm$0.13	1.26$\pm$0.25
CO2.2	16.65$\pm$0.38	17.53$\pm$0.66	8.39$\pm$0.74	19.84$\pm$0.79	13.78$\pm$0.50	14.18$\pm$0.60	17.74$\pm$0.98
CO2.3a	19.00$\pm$0.74	6.86$\pm$0.90	4.91$\pm$0.74	24.38$\pm$1.20	8.70$\pm$0.32	14.02$\pm$0.90	25.90$\pm$0.83
CO2.3b	17.96$\pm$0.93	8.97$\pm$1.09	6.80$\pm$1.03	27.70$\pm$1.68	–	8.52$\pm$1.12	35.37$\pm$1.07

Equations

$$[MgFe{]}^{\prime }\equiv \sqrt{Mgb(0.72\times Fe5270+028\times Fe5335)}$$

Full text

††pubyear: 2018††pagerange: Optical/NIR stellar absorption and emission-line indices from luminous infrared galaxies–10

1 Introduction

One challenge in modern astrophysics is to understand galaxy formation and evolution. Both processes are strongly related to the star-formation history (SFH) of galaxies. Thus, the detailed study of the different stellar populations found in galaxies is one of the most promising ways to shed some light on their evolutionary histories. So far, stellar population studies have been concentrated mainly in the optical spectral range (e.g. Bica, 1988; Worthey et al., 1994; Trager et al., 2000; Sánchez-Blázquez et al., 2006; González Delgado et al., 2015; Goddard et al., 2017; Martín-Navarro et al., 2018). In the near infrared, (0.8-2.4$\mu$m, NIR) even with some work dating back to the 1980s (e.g. Rieke et al., 1980), stellar population studies have just started to become more common in the last two decades (Origlia et al., 1993; Origlia et al., 1997; Riffel et al., 2007, 2008, 2009; Cesetti et al., 2009; Lyubenova et al., 2010; Chies-Santos et al., 2011a, b; Riffel et al., 2011c; Kotilainen et al., 2012; Martins et al., 2013b; La Barbera et al., 2013; Zibetti et al., 2013; Noël et al., 2013; Dametto et al., 2014; Riffel et al., 2015; Baldwin et al., 2017; Dametto et al., 2019; Alton et al., 2018; Dahmer-Hahn et al., 2018; Dahmer-Hahn et al., 2019; Francois et al., 2018, for example). Models have shown that the NIR spectral features provide very important insights, particularly into the stellar populations dominated by cold stars (e.g. Maraston, 2005; Riffel et al., 2007; van Dokkum & Conroy, 2012; Conroy & van Dokkum, 2012; Zibetti et al., 2013; Riffel et al., 2015; Röck, 2015; Röck et al., 2016). For example, the stars in the thermally pulsing asymptotic giant branch (TP-AGB) phase may be responsible for nearly half of the luminosity in the K band for stellar populations with an ago of $\sim$ 1 Gyr (Maraston, 1998, 2005; Salaris et al., 2014).

One common technique to study the unresolved stellar content of galaxies is the fitting of a combination of simple stellar populations (SSPs) to obtain the SFH. However, due to difficulties in theoretical treatment (Maraston, 2005; Marigo et al., 2008; Noël et al., 2013) and the lack of complete empirical stellar libraries in the NIR, (Lançon et al., 2001; Chen et al., 2014; Riffel et al., 2015; Villaume et al., 2017) the available SSP models produce discrepant results (e.g. Baldwin et al., 2017), thus making it very difficult to reliably analyse the SFH in the NIR.

On the other hand, the stellar content and chemical composition of the unresolved stellar populations of galaxies can also be obtained by the study of the observed absorption features present in their integrated spectra. So far, we still lack a comprehensive NIR dataset to compare with model predictions, required to make improvements to the models and to lead to a better understanding of the role played by the cooler stellar populations in the integrated spectra of galaxies.

Among the best natural laboratories to study these kinds of stellar content are infrared galaxies, sources that emit more energy in the infrared ($\sim$5-500$\mu$m) than at all the other wavelengths combined (Sanders & Mirabel, 1996; Sanders et al., 2003). The relevance of studying these galaxies lies particularly in the fact that they are implicated in a variety of interesting astrophysical phenomena, including the formation of quasars and elliptical galaxies (e.g. Genzel et al., 2001; Veilleux, 2006; Wang et al., 2006). When studying luminous infrared galaxies in the Local Universe, it is possible to obtain high-angular-resolution observations of these objects, thus allowing the investigation of their very central regions. Comparison of such objects with those at higher redshifts may help to understand the SFH over cosmic times.

With the above in mind, we obtained optical and NIR spectra of a sub-sample of galaxies selected from the IRAS Revised Bright Galaxy Sample present in the Local Universe. These galaxies are believed to be experiencing massive star formation, making them suitable for studying their most important spectral features that can be used as proxies to test and constrain stellar-population models. As part of a series of papers aimed at studying the stellar population and gas emission features, here we provide measurements for the most conspicuous emission and absorption features, and present new correlations between absorption features. The outline of the paper is as follows. In §2 we describe the observations and data reduction. The results are presented and discussed in §3. Final remarks are made in §5.

2 Observations and data reduction

Our sample is composed of 16 Local Universe (v$\rm{}_{r}\lesssim$ 6400 km s*-1*) galaxies that are very bright in the infrared (see Tab. 1). They were selected from the IRAS Revised Bright Galaxy Sample, which is regarded as a statistically complete sample of 629 galaxies, with 60 $\mu$m flux density $\gtrsim$ 5.24 Jy. Galaxies chosen for this study were those with log($L_{fir}/L_{\odot}$) $\gtrsim$ 10.10, accessible from the Infrared Telescope Facility (IRTF) and the Wyoming Infrared Observatory (WIRO, see below), and bright enough to reach a S/N $\sim$ 100 in the K-band within a reasonable on-source integration time.

2.1 Near Infrared Data

Cross-dispersed near-infrared (NIR) spectra in the range 0.8-2.4 $\mu$m were obtained on October 4, 6, and 7 in 2010 with the SpeX spectrograph (Rayner et al., 2003) attached to the NASA 3 m IRTF telescope at the Mauna Kea observing site. The detector is a 1024$\times$1024 ALADDIN 3 InSb array with a spatial scale of 0.15*′′/pixel. A 0.8${}^{\prime\prime}\times$15′′* slit was used during the observations, giving a spectral resolution of R$\sim$ 1000 (or $\sigma=$ 127 $\rm km\,s^{-1}$). Both the arc lamp spectra and the night-sky spectra are consistent with this value (Riffel et al., 2013a). The observations were done by nodding in an Object-Sky-Object pattern with typical individual integration times of 120 s and total on-source integration times between 18 and 58 minutes. During the observations, A0 V stars were observed near each target to provide telluric standards at similar air masses. These stars were also used to flux calibrate the galaxy spectra by using black body functions to calibrate the observed spectra of the standard stars. The seeing varied between 0.4*′′–0.7′′* over the different nights and there were no obvious clouds.

We reduced the NIR observations following the standard data reduction procedures given by Riffel et al. (2006); Riffel et al. (2013b). In short, spectral extraction and wavelength calibration were performed using spextool, software developed and provided by the SpeX team for the IRTF community (Cushing et al., 2004). The area of the integrated region is listed in Tab. 1. Each extraction was centred at the peak of the continuum-light distribution for every object of the sample. No effort was made to extract spectra at positions different from the nuclear region, even though some objects show evidence of extended emission, as this goes beyond the scope of this analysis. Telluric absorption correction and flux calibration were applied to the individual 1-D spectra by means of the IDL routine xtellcor (Vacca et al., 2003).

2.2 Optical Data

For completeness, the same sample was also observed in the optical range on nearly the same dates as the NIR data were collected with the WIRO Long Slit Spectrograph. The instrument is attached to the University of Wyoming’s 2.3-meter telescope, located on Jelm Mountain at WIRO. The Cassegrain-mounted instrument uses a Marconi 2k$\times$2k CCD detector. During our observations we used a 900 l/mm grating in first order to obtain spectra from approximately 4000-7000 Å calibrated with a CuAr comparison lamp. Given our 4-arc-second slit oriented North-South, the resolution was R$\sim$1200. Due to the relatively large spatial extent of these low-redshift objects, we offset the telescope pointing by two arc-minutes to obtain sky spectra uncontaminated by galaxy light. The seeing varied between 1-2 arc-seconds during the nights of observation. We reduced the spectra using standard techniques in IRAF111IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.. Table 1 shows the observation log along with extraction apertures. The 1-D wavelength and flux-calibrated spectra were then corrected for redshift, determined from the average $z$ measured from the position of [S iii] 0.953$\mu$m, Pa$\delta$, He i 1.083$\mu$m, Pa$\beta$ and Br$\gamma$.

Examples of the final reduced spectra, from optical to NIR ($\sim$0.4$\mu$m —2.4$\mu$m) are presented in Figure 1, for the remaining galaxies see Appendix A . For each galaxy we show the optical, $z$+$J$, $H$ and $K$ bands, from top to bottom, respectively. It is worth mentioning that the optical and NIR data do not share the same apertures, and the slit was not generally oriented at the same position angles. However, since we are interested in the nuclear region, the different slit orientations should not introduce large discrepancies in the measurements. The ordinate axis represents the monochromatic flux in units of $\rm 10^{-15}\,erg\,cm^{-2}\,s^{-1}$ Å*-1*. The position of the most common and expected emission and absorption lines are indicated as dotted (red) and dashed (blue) lines, respectively.

3 results

3.1 Emission-line spectra

A visual inspection of the data reveals a wide diversity of emission-line strengths and species. The most common emission features detected are: H$\beta$, $[$O iii] 4959, 5007 Å, $[$N ii] 6548, 6583 Å, H$\alpha$, $[$S ii] 6716,6730 Å, $[$S iii] 9531 Å, Pa$\delta$, [C i] 9824, 9850 Å, Pa$\beta$, He i 10830 Å, [P ii] 11886 Å, [Fe ii] 12570, 16436 Å, Pa $\alpha$ H2 19570 Å, H2 21218 Å, and Br$\gamma$.

Emission-line fluxes for each object of the sample were measured by fitting a Gaussian function to the observed profile and then integrating the flux under the curve. The LINER software (Pogge & Owen, 1993) was used for this purpose. No attempt to correct for stellar absorption was made before measuring the emission lines. This was done because NIR models with adequate spectral resolution (to allow the measurements of the weaker emission lines) are not available for the younger ages. Martins et al. (2013a) have shown that the underlying stellar population has only a strong effect on the hydrogen recombination emission lines, with the largest differences in fluxes being about 25 per cent. This value is within the largest uncertainties on the fluxes values too. For completeness, we have not subtracted the stellar features from the optical range too.

The results, including 3$\sigma$ uncertainties, are listed in Tables 3 and 4. For most of our targets, these measurements are made public for the first time. In addition, we computed the extinction coefficient, Cext, for the NIR using the intrinsic value of 5.88 for the flux ratio of Pa$\beta$/Br$\gamma$ (Hummer & Storey, 1987, using case B). The Cardelli et al. (1989) extinction law was used, and the values obtained for the coefficients are listed in Tables 3 and 4.

3.2 The continuum spectra

The main goal of this section is to characterize the continuum emission observed in our sample and compare it to other data in the literature. To help in the visual inspection222 Emission lines and equivalent widths of the absorption features were measured on the spectra previous to normalization. of the individual spectra, we normalized the continuum emission to unity in two regions free from emission/absorption features taken from Riffel et al. (2011b). The NIR spectra were normalized at 20925 Å and then sorted according to their continuum shapes. For a proper comparison with the optical portion of the spectrum, we normalized the optical spectra at 5300 Å and plotted them in the same order as the NIR spectra (Figs. 5 and 5).

A first-order inspection of Figure 5 allows us to infer that, contrary to what happens in Seyfert galaxies (Riffel et al., 2006), there seems to be no correlation between activity type (LINERs or SFGs) and continuum shape. In fact, these very bright infrared galaxies present a continuum shape very similar to what is found in fainter Hii sources and normal galaxies, as reported by Martins et al. (2013a), which may indicate that the LINER spectrum of these galaxies is powered by starburst instead of a low-luminosity AGN. In addition, the continua of all the optical spectra look very similar.

A large diversity of atomic absorption lines and molecular bands is also apparent in the spectra. These features are seen from the very blue optical end to the red end of the observed NIR spectral region. The most common atomic absorption features are due to Ca i, Ca ii, Fe i, Si i, Na i and Mg i, besides the prominent absorption bands of CH, MgH, TiO, VO, ZrO and CO. These features are identified in Fig. 1. It is clear in these figures that some of the most important features predicted for intermediate-age stellar populations, which are expected to be enhanced in the RGB and TP-AGB stellar phases (Maraston, 2005; Riffel et al., 2007, 2015), are detected in the spectra. Among these features are the ZrO/CN/VO at 9350 Å, the 10560 Å VO, 1.1 $\mu$m CN and 1.6 $\mu$m and 2.3 $\mu$m  CO bands.

3.2.1 Towards a homogeneous NIR index definition

The Equivalent Widths (EWs) of these features offer coarse but robust information about the stellar content of a galaxy spectrum, and therefore they can be used as powerful diagnostics of the stellar content of galaxies. Contrary to the optical range, where there exist indices defined in a homogeneous way by the Lick group (see Worthey et al., 1994, and references), in the NIR there is no such homogeneous set of definitions covering the full NIR wavelength range, and authors tend to use their own definitions (e.g. Riffel et al., 2007, 2008; Silva et al., 2008; Mármol-Queraltó et al., 2009; Cesetti et al., 2009; Kotilainen et al., 2012; Riffel et al., 2011a, 2015; Röck et al., 2017), and therefore it is very difficult to compare results from different investigations.

With this in mind, here we create a set of definitions for absorption features found in the NIR. We used two SSPs from the IRTF-based emiles models (Vazdekis et al., 2016; Röck, 2015; Röck et al., 2016), with 1.0 Gyr and 10 Gyr, solar metallicity, calculated with the PADOVA evolutionary tracks and with \sigma=228$$\rm km\,s^{-1}. We added up their light fractions (normalized to unity at $\lambda$=12230Å) as follows:

$F_{\rm comb}=0.5\frac{F_{\lambda}^{1Gyr}}{F_{\lambda=12230}^{1Gyr}}+0.5\frac{F_{\lambda}^{10Gyr}}{F_{\lambda=12230}^{10Gyr}}$.

To this resulting spectrum we added Gaussians to model emission-lines profiles. These lines are located at the wavelengths of the most common emission lines detected in galaxies in this spectral region (see Sec. 3.1) with Full Width at Half Maximums (FWHMs) characteristic of galaxies observed with SpeX with the configuration used here (25Å $\rm~{}\lesssim~{}FWHM~{}\lesssim$ 40Å) with arbitrary flux values. We employed the elprofile routine of the ifscube package333available at: https://bitbucket.org/danielrd6/ifscube.git (Ruschel-Dutra, in preparation). Using this simulated spectrum we defined the line limits and continuum band passes as illustrated in Fig. 3 and listed in Table 5.

We have measured the EWs for the most prominent absorption features using an updated python version of the pacce code (Riffel & Borges Vale, 2011). In this code version, the EW uncertainties are assumed to be the standard deviation of 1000 EWs measurements of simulated spectra created by perturbing each flux point by its uncertainty through a Monte Carlo approach. The line definitions used are listed in Table 5, and the measured values are in Tables 6 and 7. In order to have a sample of early type galaxies (ETG) to compare our results with, we have collected NIR spectra from the literature and measured the EW of the absorption features with the same definitions used for our sample. Tables 8 and 9 present the measurements for the sample of galaxies presented in Baldwin et al. (2017). For four of the galaxies we were able to find Sloan Digital Sky Survey data used to measure the optical EW, while for the remaining objects we collected the values of Fe5015, Mgb and Fe5270 from McDermid et al. (2015). We also measured the values from the spectra presented by Dahmer-Hahn et al. (2018), which values are listed in Tab. 10.

4 discussion

4.1 Emission Lines

In order to compare the frequency of occurrence of the emission lines in our sample with what is seen in Seyfert galaxies, we show a histogram in Fig. 2 where the lines found here are compared to those of Riffel et al. (2006). What clearly emerges from this figure is that $[$S iii], He i, and Pa$\beta$ lines are less frequent in our sample (occurring in $\sim$60% of the sources) than in Seyferts (present in almost all of the objects). On the other hand, we find a higher frequency of occurrence of lines of [C i] ($\sim$ 65%), [P ii] ($\sim$ 40%) and [Fe ii] ($\sim$ 65%) than in Sy 1 objects, and a similar rate as in Sy 2s. The remaining emission lines occur with similar frequencies in the present sample and in Seyferts (see also Lamperti et al., 2017). Lines that are less frequent in the present sample compared to AGNs are located in regions with strong stellar features. Thus, it is possible that the absence of these features is because they are intrinsically weaker than in AGNs and/or diluted by the broad absorption features that dominate the $z+J$ band. Note though that for three objects (NGC 1055, NGC 6835 and NGC 520, see Fig 5), our spectral range excludes the $[$S iii], He i, and [C i] emission lines. If present in these spectra, they would show up in $\sim$80% of our sample.

It is worth mentioning that the kinematics of the [S iii], [Fe ii]  and H2 lines as well as the excitation mechanisms of the [Fe ii] and H2 lines of the galaxies of this sample were explored in Riffel et al. (2013b). However, the low-ionization forbidden lines of [C i] (i.e. $\lambda$ 9850 Å) and [P ii] (i.e. $\lambda$ 11886 Å), also detected in our sample, were not yet analyzed. Although the [P ii] line is stronger compared to [Fe ii] $\lambda$ 12570 Åin Sy 2s (Riffel et al., 2006) than in the other types of galaxies, the detection of [P ii] lines is surprising here. This is because at Solar metallicity, Phosphorus is about 1000 times less abundant than Carbon (Ferguson et al., 1997) and 100 times less abundant than Iron (Oliva et al., 2001). Hence, if the P/C abundance is near to solar, the [P ii] lines should not be present, unless other strong abundant elements are much more optically thick than they appear. A similar problem is found in some quasars for which broad absorption lines of P v $\lambda\lambda$1118,1128 Å are detected and extreme abundances ratios for P/C are found (Hamann, 1998; Hamann et al., 2001; Borguet et al., 2012). According to Oliva et al. (2001) for a solar Fe/P $\sim$ 100 abundance ratio, one expects that $\frac{[\text{Fe\,{ii}}]}{[\text{P\,{ii}}]}=50$, similar to what is expected for supernova remnants. The NIR [P ii] emission lines may probably help to set some constraints on the abundance of Phosphorus in galaxies.

As discussed in Oliva et al. (2001) bright [Fe ii] lines can only be formed in regions where hydrogen is partially ionized. Such regions of hot, partially ionized gas can only be produced in an efficient way by shocks and/or photoionisation by soft X-rays. According to these authors, [Fe ii]/[P ii] can be used to distinguish between shocks (ratio $\gtrsim$20) and photoionisation (ratio $\lesssim$ 2). In order to test this hypothesis, we plotted in Fig. 6 [Fe ii]/[P ii] $\times$ [C i]/[P ii] for our sample as well as the Seyfert galaxies of Riffel et al. (2006). As can be seen in this figure, there is a good correlation and no clear separation between the SFGs and the Seyferts, suggesting that the dominant excitation mechanism is the same for the three ions. Furthermore, due to the low values derived for the [Fe ii]/[P ii] ratio, that excitation mechanism might be expected to be photoionisation based on the arguments of Oliva et al. (2001). To test this, we have computed photoionisation models using cloudy/C17.01444Available at https://www.nublado.org. (Ferland et al., 2017) updated with the next release of collisional strengths for [P II] (taken from Tayal, 2004) as well as with new transition probabilities555They are a combination of data taken from the MCHF/MCDHF Database at http://nlte.nist.gov/MCHF/ and data from the NIST Atomic Spectra Database at https://www.nist.gov/pml/atomic-spectra-database (private communication), however these models are not able to reproduce the observed line ratios, underestimating both (the models values for both ratios are nearly zero). This may be due to the fact that these lines are not in fact excited by photoionisation, but mostly driven by shocks.

4.2 Absorption Features

It is crucial to be able to derive ages and chemical composition in order to understand the dominant underlying unresolved stellar content of galaxies (Röck et al., 2017). So far, the NIR is lacking a clear procedure based on absorption-line strengths. The obvious choice to do this kind of study is using stellar clusters as probes, instead of the use of more complex star-forming objects. However, while observations of the integrated spectra of stellar clusters in the optical region have been available for almost 30 years (e.g. Bica, 1988) in the NIR such observations are very difficult since the light emitted by the stars of the clusters in the NIR bands is dominated by a few very bright stellar phases making it difficult to get reliable integrated spectra of such objects in the NIR (e.g Lyubenova et al., 2010; Riffel et al., 2011c).

In order to have a more homogeneous data-set, in addition to the data-set we present here representing complex SFHs of SFGs (§ 2), we collected spectra of nearby ETGs (which tend to have less complex SFHs than our sample) observed similarly as those in the present work. Our final data set representing the older stellar population is composed of 12 ETG selected in order to span a wide range of ages (1-15 Gyr) at approximately solar metallicity and observed by Baldwin et al. (2017) using Gemini/GNIRS in the cross-dispersed mode ($\sim 0.8-2.5\mu m;R\sim 1700$; $\sigma\sim$75 $\rm km\,s^{-1}$) plus 6 ETG selected from the Calar Alto Legacy Integral Field Area Survey (CALIFA Sánchez et al., 2016) and observed by Dahmer-Hahn et al. (2018) using the TripleSpec spectrograph attached to the Astrophysical Research Consortium (ARC) 3.5-meter telescope ($\sim 0.95-2.45\mu m;R\sim 2000$; $\sigma\sim$64 $\rm km\,s^{-1}$). In addition to these NIR spectra, we also collected, when available, the optical spectra of the sources. In the case of Baldwin et al. (2017) galaxies, the optical spectra where taken from the Sloan Digital Sky Survey (Ahn et al., 2014), while for the sample of Dahmer-Hahn et al. (2018) we took the data from the (CALIFA Sánchez et al., 2016). The optical and NIR indices were measured by us using the definitions of Tab. 5 and are listed as online material in Tabs. 8, 9 and 10.

4.2.1 Previous NIR index - index correlations

Due to the lack of adequate data sets to test predictions of NIR data, compared to the optical (see Thomas et al., 2003, for example), there are only a few studies trying to understand the behaviour of NIR $\times$ NIR indices. For instance, Mármol-Queraltó et al. (2009) studied a sample of early type galaxies and found a strong correlation between $C_{2}$4668 and NaI2.20 indices. In Fig. 7a we show the Mármol-Queraltó et al. (2009) measurements (open diamonds) and the literature compilation presented by Röck et al. (plus symbols, 2017) together with our data (squares). Even though we only measured both indices for 4 sources, this correlation seems to still hold for SFGs, which populate the lower left end of the correlation (Fig. 7a).

Using a similar approach, Cesetti et al. (2009) reported a trend of correlation of the optical Mg2 band with NIR indexes, such as NaI2.20, CaI2.26 and CO2.2 for early type galaxies. In Fig 7b,c and d we plotted our sample (filled squares), together with those of Cesetti et al. (2009, open diamonds) and Kotilainen et al. (open triangles 2012) for early type sources. Additionally we also added the inactive spirals (octagons, LTG-K12) of Kotilainen et al. (2012). From Fig. 7d we have excluded the two Seyfert galaxies (NGC 660 and NGC 6814) since the CO band can be very diluted in these kind of sources (Riffel et al., 2009; Burtscher et al., 2015).

From Fig. 7 it is clear that the trend seems to hold for NaI2.20 $\times$ Mg2, while for CaI2.26 $\times$ Mg2 there is no clear correlation, and in the case of CO2.2 $\times$ Mg2 instead of a positive correlation there seems to be an inverse correlation. Additionally there seems to be a segregation between early and late-type galaxies in this plot (panel d). This indicates that CO is enhanced in younger stellar populations, in agreement with the predictions of the Maraston (2005) models as shown in Riffel et al. (2007).

To help in the interpretation of these results, on these index-index diagrams we have over-plotted the new optical-to-NIR IRTF-based stellar population synthesis models of the E-MILES team (Vazdekis et al., 2012, 2016; Röck et al., 2016). The models employed are those computed using the PADOVA isochrones (Girardi et al., 2000), with ages in the range 0.3 Gyr $<t<$ 15.0 Gyr and metallicities within [Fe/H] = -0.40, [Fe/H] = 0.00 and [Fe/H] = 0.22 with two different spectral resolutions ($\sigma$=60$\rm km\,s^{-1}$ and $\sigma$=228$\rm km\,s^{-1}$, the shaded area represents the differences caused by $\sigma$). We also plotted TP-AGB heavy (see Zibetti et al., 2013, for a comparison between TP-AGB heavy and light models), Pickles-based models of Maraston & Strömbäck (2011, M11 hereafter), which do have the same prescription than Maraston (2005) models but with a higher spectral resolution (R = 500) than the 2005 models, therefore making them more suitable for our comparisons. However, it is important to have in mind that M11 models do have a poorer spectral resolution than our data, the effects on the indices strengths by degrading the resolution to M11 models is within the uncertainties of our measurements. These models are shown as open brown stars and are only available for solar metalicity. What emerges from this exercise is that the models in general are not able to predict the NIR indices and that there is a segregation between early (open diamonds and plus markers) and late-type (filled squares and octagons) galaxies in these diagrams. The upper panels show significantly larger NaI2.20 index values than predicted by the models with standard IMF. Both the optical Mg and C dominated indices are stronger than the models for the most massive galaxies (i.e. the ones with the largest index values). In the case of the NaI2.20 index, Röck et al. (2017) concluded that for early-type sources the large values obtained for this index are due to a combination of a bottom-heavy initial mass function and the [Na/Fe] abundances. On the other hand, Alton et al. (2018) found that their sample of massive ETGs is consistent with having a Milky Way-like IMF, or at most a modestly bottom-heavy IMF, and suggested that their extreme abundance values for Na, in the cores of massive ETGs, may be explained by the metallicity-dependent nucleosynthetic yield of Na.

The lower panels of Fig. 7 show that the ETGs are in better agreement with the predicted values. However, about half of our SFGs sample show stronger CaI2.26 and CO2.2 values than predicted by the models. From these plots, we also can infer that the TP-AGB phase does not change substantially the CO index, once the solar metalicity M11 models are in agreement with the E-MILES ones for the younger ages ($t\lesssim$ 1 Gyr), with a large discrepancy for the older ages. Besides age, metallicity appears as an additional discriminator for the measured strengths of CO bands, with low metalicity ([Fe/H] = -0.40) and intermediate ages ($\sim$350Myr) showing the largest values for the CO2.2 index. This is in agreement with the previous findings of Kotilainen et al. (2012), who found that the evolved red stars completely dominate the NIR spectra, and that in this age range, the hot, young stars contribution to the EWs is virtually nonexistent. So far, to fully access these younger stellar content of the galaxies it is necessary to fit the full spectrum, taking the continuum into account (see Baldwin et al., 2017; Dahmer-Hahn et al., 2018, for example). However, this is beyond the scope of the present paper and will be the subject of a future investigation (Riffel et al., in preparation). On the other hand, the lower values of the CO index presented by the ETGs are also not explained by the models, with M11 models underestimating and E-MILES models overestimating the main locus occupied by these sources.

4.2.2 New index - index correlations

Because we measured a large set of lines for our sample, we have tried to find new correlations among the different absorption features by plotting all the EWs listed in Tab. 6 and 7, as well as literature data (Tabs. 8 to 10) against each other. From these, we removed the correlations already discussed above (Fig. 7) as well as the optical $\times$ optical indices correlations since these are well studied 666The NaI2.20 and CO2.2 are well studied, however, we decided to keep them here for diagrams distinct from those presented in Fig. 7 because correlations with other lines may help to shed some light in the understanding of the mechanisms driving these lines.. Since the CaT lines are correlated (e.g. Cenarro et al., 2001), we only used CaT2 in our search for correlations. The final set of optical versus NIR and NIR versus NIR indices the correlations are shown in Figs. 8 and 9, together with a linear regression using the orthogonal distance regression (ODR) method that takes errors both in the x and y variables into account (Boggs & Rogers, 1990). We note that when it was not possible to measure one of the indices used in the correlations, we have removed the galaxy from the plots and regression. In addition, we only considered the cases were both indices were measured at least for 6 sources. To help to understand these plots we have over-plotted the same model set as discussed above.

What emerges from Fig. 8 is that both model sets are able to predict well all the measured values for the optical indices. In the NIR, however, the models fail in their predictions, except for CO2.2 and ZrO, with E-MILES making better predictions of strengths than M11, especially in the case of atomic absorption features. In addition, there is a clear separation of the ETGs and SFGs on the G4300$\times$MgI1.7, G4300$\times$NaI2.20, Fe4531$\times$MgI1.7 and Mg${}_{1}\times$NaI2.20 diagrams with ETGs in general showing higher values for both optical and NIR indices. A less evident separation of ETGs and SFGs is observed on the G4300$\times$CO2.2 and Mg${}_{b}\times$NaI2.20 diagrams, while no separation is observed for the Fe5782$\times$ZrO and TiO${}_{1}\times$ CO2.2 diagrams.

The optical indices (G4300, Fe4531 and Mg1) are not very sensitive to the $\alpha$/Fe ratio while G4300 is mainly sensitive to the C and O abundances (Thomas et al., 2003). This may indicate that the MgI1.7 and NaI2.20 indices are also sensitive to C and/or O abundances. This is also in agreement with the findings of Röck et al. (2017) who suggested that [C/Fe] enhancement might contribute to the values observed for NaI2.20 in ETGs. However, the good correlation of NaI with Mgb may also indicate that this index is $\alpha/Fe$ dependent, since Mgb is sensitive to changes in the $\alpha/Fe$ ratio (Thomas et al., 2003). red The CO2.2 index values are well described by the model predictions for both SFGs and ETGs, with an age-metallicity dependence for the SFGs and no evidence of strong changes on their strengths caused by the amount of TP-AGB stars (see above). This is additionally supported by the CO2.2 $\times$ TiO1 diagram, where M11 models, independent of age, do populate the locus filled by the ETGs, while E-MILES models do not reproduce the larger TiO and smallest CO strengths. The CO and TiO1 correlation is not unexpected since these absorptions depend on O being available. The models do show that ZrO is more metallicity dependent while TiO1 seems to be age dependent. In addition, some ETG show TiO1 values larger than the models (specially E-MILES models), which can be interpreted as an IMF effect (see La Barbera et al., 2013). In the case of the Mg-dominated indices (in the NIR and optical) the large values for these indices can be associated with the most massive ETGs, and can be explained by an [Mg/Fe] enhancement (e.g. Worthey et al., 1992; Martín-Navarro et al., 2018).

The correlations found from this exercise for the NIR indices are shown in Fig. 9. One particularly relevant correlation is CO1.6b$\times$CN11, as the CN11 index is believed to be heavily dominated by the AGB evolutionary phase and particularly by C stars (Maraston, 2005). Almost 50% of our SFG do show CN11$\gtrsim$10Å, with a mean value $\sim$20% larger than in ETG (see Fig. 13) and are consistent with the intermediate age (0.3 - 2 Gyr) models. M11 models do cover better the space of values of the measurements, but all the older ages M11 models (t$\gtrsim$ 3 Gyr) do predict more or less constant values for CN11 (the same hapens for CO1.6b). The ETGs are more or less matched by SSP models with old ages and no indication of an intermediate age population is required to explain the absorption features of these sources, once, their strengths in some cases are smaller than those of the older E-MILES SSPs. The CO2.2 and CO2.3a,b (also CO1.5a and CO1.5b) indices are to some extent described by the models, with larger values predicted for intermediate age SSPs.The remaining strengths are not predicted by the models and no clear separation is found for SFGs and ETGs.

With the aim of understanding the behaviour of the NIR indices, we plotted them against the [MgFe]′ index of Thomas et al. (2003) defined as:

[TABLE]

which, for this sample with a small range in metallicity, is basically an age-indicator and is completely independent of the $\alpha/Fe$ ratio. Assuming that the ETGs used here are objects with relatively normal old stellar populations, which is a valid assumption since their full spectra can be well fitted with SSP models (see Baldwin et al., 2017; Dahmer-Hahn et al., 2018, for details). This can also be seen in Figs. 11 to 12, where the ETGs do in general show less scatter in the [MgFe]′ index than the SFGs. This indicates a more complex SFH for the latter, most likely with a strong contribution from intermediate ($\sim$1 Gyr) stellar populations. In order to test the effect of a more complex SFH on the NIR strengths, we show in Fig. 13 histograms comparing the strenght distributions between SFG and ETG. Except for a few indices (ZrO, MgI1.48, MgI1.50, CO1.5a, FeI1.58, CO1.5c, MgI1.7, NaI2.20), the mean value for SFG is larger than that for ETG. This more complex SFH can also explain why the CN and CO bands are in general stronger for the SFGs than the ETGs. These bands are enhanced by the short-lived younger red giant branch (RGB) and thermally pulsing asymptotic giant branch (TP-AGB) stars (Maraston, 2005; Riffel et al., 2007, 2015). According to Maraston (1998), these stars can be responsible for up to 70 per cent of the total flux in the NIR. However, for the case of the NaI2.20 index, (Röck et al., 2017) constructed models using enhanced contribution from AGB stars and found that these stars have only a very limited effect on the model predictions and do not improve significantly the fit of the model NaI2.20 indices. They also show that small fractions (3 per cent) do have a similar impact on NaI2.20 than those with larger amounts of these stars. This result is consistent with our findings that NaI2.20 index has a mean value $\sim$20% larger in ETG than in SFG.

In general, the NIR line strengths are not well reproduced by any set of models, suggesting that the SFH of the galaxies cannot be recovered when only using NIR indices. Our results are in agreement with the finding of Baldwin et al. (2017) who have studied the SFH of a sample of ETG by fitting different SSP models and found that the SFH vary dramatically among the different EPS models when fitting NIR data, with higher spectral resolution models producing more consistent results. They also found variations in ages in the NIR tend to be small, and largely encoded in the shape of the continuum. This was also noticed in Riffel et al. (2015) who suggested that TP-AGB stars contribute noticeably to a mean stacked NIR spectrum made up with mostly late type galaxies hosting a low luminosity AGN, from the Palomar survey (Mason et al., 2015). This result was obtained by fitting a mix of individual IRTF stars to the mean galaxy spectrum. Nevertheless, in this same work we have shown that other evolved stars (red giants, C-R and E-AGB stars) can reproduce most of the absorption features detected, without having to resort to stars in the TP-AGB phase.

5 Final Remarks

We analysed long-slit spectra spanning optical to near-infrared wavelengths of 16 infrared-luminous star-forming galaxies with the aim of offering the community a set of emission and absorption feature measurements that can be used to test the predictions of the forthcoming generations of stellar population models. The optical and NIR spectra were obtained at WIRO and at SpeX/IRTF, respectively. In addition to these, we collected literature spectra of early-type galaxies and performed the equivalent width measurements using a new homogeneous set of continuum and band pass definitions. The main findings can be summarized as follows:

• •

All our sources display H2 emission, characteristic of the star-forming nature of our sample. In the optical they clearly display H i emission lines. However, NGC 1055 and NGC 1134 show a NIR spectrum free of H i emission lines. We interpret this latter result as the result of the low sensitivity of the NIR detector in this wavelength interval, thus the expected Br$\gamma$ fluxes are below the detection limit.

• •

The continua are dominated by stellar absorption features. The most common features are due to Ca i, Ca ii, Fe i, Na i, Mg i, plus prominent absorption bands of: TiO, VO, ZrO and CO. In most cases (70%) ,the stellar continua also show evidence of dust extinction.

• •

We present new definitions of continuum and line band passes for the NIR absorption lines. These definitions were made taking into account the position of the most common emission lines detected in this wavelength range.

• •

We report EW measurements for 45 indices, including both optical and NIR features. We also present measurements for most of these indices in spectra of ETGs taken from literature. To the best of our knowledge, they represent the most complete set of EW measurements reported in the literature to date, and can be used to test the predictions of stellar population models from the optical to the NIR.

• •

We looked for correlations among the different absorption features, presenting as the most robust ones those with a Pearson correlation coefficient r$>$0.6. In addition to the already-known correlations in the optical region, we propose here correlations between optical and NIR indices, as well as correlations between different NIR indices, and compare them with model predictions.

• •

While for the optical absorption features the new generation of models, with scaled-solar abundance ratios and standard IMF, share the same locus as the observed data points, they fail to predict the strengths of most of the NIR indices for the SFGs, while in the case of the early-type sources they roughly reproduce the observations. This may indicate more complex SFHs for the SFGs, which we interpreted as a strong contribution from the younger stellar populations, thus explaining the fact that the CN and CO bands are in general larger for the SFGs than the ETGs. These bands are enhanced in stars in the TP-AGB phase, however, they seems to have a limited impact on the indices of ETGs.

Acknowledgements

We are grateful to the referee for insight that has improved the quality of this paper. RR and RAR thank CNPq, CAPES and FAPERGS for financial support for this project and to Luis Colina for helpful discussions on this project. The authors are also grateful to Richard McDermid and Christina Baldwin for kindly offering information on the ETGs used in this work, as well as, to Jari Kotilainen for sharing information on his sample of objects. We also thank Peter van Hoof for providing the new collisional strengths for [P ii] lines, which will be part of the new cloudy release, as well as for useful discussions. ARA thanks CNPq for financial support. RFP acknowledges financial support from the European Union’s Horizon 2020 research and innovation program under the Marie Sklodowska-Curie grant agreement No. 721463 to the SUNDIAL ITN network. AV acknowledges support from grant AYA2016-77237-C3-1-P from the Spanish Ministry of Economy and Competitiveness (MINECO).

Appendix A Final reduced spectra

Final reduced and redshift-corrected spectra for the remaining sample. Available as online material.

Appendix B Literature data

Here we present the measurements using the index definitions listed in Tab. 5 for the literature data. The data used here are those of Dahmer-Hahn et al. (2018) and Baldwin et al. (2017). For the latter we found optical Sloan Digital Sky Survey data (Ahn et al., 2014) for four sources. For the remaining objects we collected the values of Fe5015, Mgb and Fe5270 from McDermid et al. (2015), while for the sources of Dahmer-Hahn et al. (2018) the optical data were taken from the Calar Alto Legacy Integral Field Area Survey (califa Sánchez et al., 2016) and we measured the EW of the optical lines.