A 1.3 cm line survey toward Orion KL

Y. Gong; C. Henkel; S. Thorwirth; S. Spezzano; K. M. Menten; C. M. Walmsley; F. Wyrowski; R. Q. Mao; B. Klein

doi:10.1051/0004-6361/201526275

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A&A, 581 (2015) A48

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Issue		A&A Volume 581, September 2015


Article Number		A48
Number of page(s)		60
Section		Interstellar and circumstellar matter
DOI		https://doi.org/10.1051/0004-6361/201526275
Published online		02 September 2015

A&A 581, A48 (2015)

A 1.3 cm line survey toward Orion KL^⋆,^⋆⋆

Y. Gong¹^,2^,3, C. Henkel²^,4, S. Thorwirth⁵, S. Spezzano⁵, K. M. Menten², C. M. Walmsley⁶^,7, F. Wyrowski², R. Q. Mao¹ and B. Klein²^,8

¹ Purple Mountain Observatory & Key Laboratory for Radio Astronomy, Chinese Academy of Sciences, 2 West Beijing Road, 210008 Nanjing, PR China
e-mail: ygong@pmo.ac.cn
² Max-Planck Institut für Radioastronomie, Auf Dem Hügel 69, 53121 Bonn, Germany
³ University of Chinese Academy of Sciences, No. 19A Yuquan Road, 100049 Beijing, PR China
⁴ Astronomy Department, King Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia
⁵ I. Physikalisches Institut, Universität zu Köln, 50937 Köln, Germany
⁶ Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 5, 50125 Firenze, Italy
⁷ Dublin Institute of Advanced Studies, Fitzwilliam Place 31, Dublin 2, Ireland
⁸ University of Applied Sciences Bonn-Rhein-Sieg, Grantham-Allee 20, 53757 Sankt Augustin, Germany

Received: 8 April 2015
Accepted: 10 June 2015

Abstract

Context. The nearby Orion Kleinmann-Low nebula is one of the most prolific sources of molecular line emission. It has served as a benchmark for spectral line searches throughout the (sub)millimeter regime.

Aims. The main goal is to systematically study the spectral characteristics of Orion KL in the λ ~ 1.3 cm band.

Methods. We carried out a spectral line survey with the Effelsberg-100 m telescope toward Orion KL. It covers the frequency range between 17.9 GHz and 26.2 GHz, i.e., the radio “K band”. We also examined ALMA maps to address the spatial origin of molecules detected by our 1.3 cm line survey.

Results. In Orion KL, we find 261 spectral lines, yielding an average line density of about 32 spectral features per GHz above 3σ (a typical value of 3σ is 15 mJy). The identified lines include 164 radio recombination lines (RRLs) and 97 molecular lines. The RRLs, from hydrogen, helium, and carbon, stem from the ionized material of the Orion Nebula, part of which is covered by our beam. The molecular lines are assigned to 13 different molecular species including rare isotopologues. A total of 23 molecular transitions from species known to exist in Orion KL are detected for the first time in the interstellar medium. Non-metastable (J>K) ¹⁵NH₃ transitions are detected in Orion KL for the first time. Based on the velocity information of detected lines and the ALMA images, the spatial origins of molecular emission are constrained and discussed. A narrow feature is found in SO₂ (8_1,7 − 7_2,6), but not in other SO₂ transitions, possibly suggesting the presence of a maser line. Column densities and fractional abundances relative to H₂ are estimated for 12 molecules with local thermodynamic equilibrium (LTE) methods. Rotational diagrams of non-metastable ¹⁴NH₃ transitions with J = K + 1 to J = K + 4 yield different results; metastable (J = K) ¹⁵NH₃ is found to have a higher excitation temperature than non-metastable ¹⁵NH₃, also indicating that they may trace different regions. Elemental and isotopic abundance ratios are also estimated: He/H = (8.7 ± 0.7)% derived from the ratios between helium RRLs and hydrogen RRLs; ¹²C/¹³C = 63 ± 17 from ¹²CH₃OH/¹³CH₃OH; ¹⁴N/¹⁵N =100 ± 51 from ¹⁴NH₃/¹⁵NH₃; and D/H = (8.3 ± 4.5) × 10^-3 from NH₂D/NH₃. The dispersion of the He/H ratios derived from Hα/Heα pairs to Hδ/Heδ pairs is very small, which is consistent with theoretical predictions that the departure coefficients b_n factors for hydrogen and helium are nearly identical. Based on a non-LTE code that neglects excitation by the infrared radiation field and a likelihood analysis, we find that the denser regions have lower kinetic temperature, which favors an external heating of the hot core.

Key words: astrochemistry / ISM: individual objects: Orion KL / ISM: molecules / radio lines: ISM / surveys

^⋆

Tables 2 and 4 and appendices are available in electronic form at http://www.aanda.org

^⋆⋆

The reduced spectra as FITS files are only available at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/581/A48

© ESO, 2015

1. Introduction

One of the best ways to study the physical and chemical conditions of astronomical objects is via spectral line surveys. In fact, they are the only means to obtain a completely unbiased view of the molecular inventory of a given source. However, despite the long history of powerful large radio telescopes, only very few unbiased frequency surveys have been conducted in the centimeter wave regime. Spectroscopic features that are expected at centimeter wavelengths include radio recombination lines (RRLs) from hydrogen, helium, and carbon, and pure rotational transitions of heavy species, such as complex organic molecules that can be detected conveniently at centimeter wavelengths thanks to their small rotational constants. The molecular lines at centimeter wavelengths are likely to be optically thin, so analysis is easier than at millimeter wavelengths. Furthermore, the “line density” is not as high as at millimeter wavelengths, reducing the confusion level.

We now summarize previous surveys at centimeter wavelengths. Bell et al. (1993) surveyed W51 from 17.6 to 22.0 GHz with the NRAO 43 m telescope; Kalenskii et al. (2004) investigated the dark cloud TMC-1 at 4−6 and 8−10 GHz with the Arecibo 305 m telescope; Kaifu et al. (2004) observed TMC-1 from 9 to 50 GHz with the Nobeyama 45 m telescope; Gong et al. (2015) measured IRC +10216 from 17.8 to 26.3 GHz with the Effelsberg 100 m telescope finding 23 new transitions from species detected previously; and Remijan et al. (2008a) surveyed Sgr B2 near the Galactic center with the Green Bank Telescope (GBT) in the course of their GBT PRIMOS project¹ which led to detections of complex organic molecules, such as CH₂CHCHO (propenal), CH₃CH₂CHO (propanal), and CNCHO (cyanoformaldehyde; Hollis et al. 2004; Remijan et al. 2008a). They have shown that the centimeter wave range of the electromagnetic spectrum has an enormous potential for detecting new molecules and new transitions.

In this study, we present a 1.3 cm line survey toward the Orion Kleinmann-Low nebula (Orion KL). Orion KL has a high luminosity of ~10⁵L_⊙ (Wynn-Williams et al. 1984), harboring several luminous embedded IR sources, and is the closest high mass star formation region at a distance of 414 ± 7 pc (Menten et al. 2007). The chemistry of this source is particularly rich owing to the high kinetic temperature of the gas caused by the evaporation of dust grain mantles and the interaction of the newly formed stars with their environment (e.g., Herbst & van Dishoeck 2009; Charnley 1997). It is thus one of the best sources for studying the spectral characteristics associated with high-mass star-forming regions.

So far, many spectral line surveys have been carried out toward Orion KL. The existing surveys, summarized in Table 1, cover a wide frequency range from 34.25 GHz to 6.8 THz. These efforts have resulted in the discovery of more than 40 molecules in Orion KL, recently including the unusual long complex molecule methyl acetate, CH₃COOCH₃ (Tercero et al. 2013), as well as a deuterated ammonium cation, NH ${}_{3}D^{+}$ $\hbox{$_{3}{\rm D}^{+}$}$ (Cernicharo et al. 2013). Meanwhile, no comprehensive view exists of our frequency range covered here (17.9–26.2 GHz), although numerous studies have targeted specific lines that it contains, such as from NH₃ or CH₃OH. Here, we therefore systematically study the λ ~ 1.3 cm (K band) characteristics of Orion KL.

Table 1

Existing line surveys of Orion KL.

2. Observation and data reduction

2.1. Observations with the Effelsberg-100 m telescope

The measurements were carried out in a position-switching mode with the primary focus λ = 1.3 cm K-band receiver of the 100-m telescope at Effelsberg/Germany², during 2012 January and April and 2013 January, March, and May. The telescope was pointed at α_J2000 = 05^h35^m14.17^s, δ_J2000 = − 05°22′46.5′′, about 12′′ south of the dust peak (see Fig. 1a). On- and off-source integration times were two minutes per scan. The newly installed Fast Fourier Transform Spectrometer (FFTS) was used as backend. Each of the two orthogonal linear polarizations was covered with a bandwidth of 2 GHz, providing 32 768 channels and resulting in a channel spacing of 61 kHz, which is equivalent to 0.7 km s^-1 at 26 GHz. The actual frequency and velocity resolutions are coarser by a factor of 1.16 (Klein et al. 2012). Several frequency setups were tuned to cover the entire frequency range from 17.9 GHz to 26.2 GHz with an overlap of at least 100 MHz between two adjacent setups. Across the whole frequency range, the FWHM beamsize varies from 35′′ to 50′′ (~40′′ at 23 GHz). The survey encompasses a total of ~25 observing hours. The focus was checked every few hours, in particular after sunrise and sunset. Pointing was obtained every hour toward the nearby pointing sources PKS 0420−01 or 3C 161 and was found to be accurate to about 5′′. Strong continuum sources (mostly NGC 7027 and 3C 286) were used to calibrate the spectral line flux, assuming standard flux densities (5.5 Jy for NGC 7027 and 2.5 Jy for 3C 286 at 22 GHz; Ott et al. 1994). The typical rms noise is about 3 − 15 mJy in 61 kHz wide channels. The conversion factor from Jy on a flux density scale (S_ν) to K on a main beam brightness temperature scale (T_mb) is T_mb/S_ν ~ 1.7 K / Jy at 18.5 GHz, 1.5 K / Jy at 22 GHz, and 1.4 K / Jy at 23.7 GHz. Local standard of rest (LSR) velocities are used throughout the paper.

Fig. 1

a) SCUBA–850 μm dust emission (contours) overlaid on the 6 cm VLA continuum image. The contours correspond to 5, 10, 20, 40, 80, 160 Jy beam^-1. The red circle represents the Effelsberg FWHM beamsize (40′′). The red box represents the mapped region of Fig. 1b. b) Continuum map of Orion KL at 230 GHz from the ALMA-SV line survey modified from Fig. 8 of Crockett et al. (2014). The contour levels are 10%, 20%, 40%, 60%, and 80% of the peak intensity of 1.406 Jy beam^-1. The positions of source n, source I, SMA1 and the BN object are indicated by the purple squares. The position of the hot core south (HC(S)) is indicated by the blue circle. The positions of infrared sources IRc1, IRc2, IRc3, IRc4, IRc5, IRc6, and IRc7 are labeled by red crosses. Blue crosses indicate millimeter wave continuum sources. The approximate positions of the compact ridge and the HC are indicated by the red dashed circles and the acronyms CR and HC. (these sources are described in more details in Sect. 3.) The physical scale is given by the horizontal bar. The synthesized beam is shown in the lower left of the panel. The reference point corresponds to (α_J2000, δ_J2000) = (05^h35^m14.350^s, −05°22′35.00′′).

For the data reduction, the GILDAS³ software package that includes CLASS and GreG was employed. During the data reduction, some inevitable defects were found at the edges of spectra, so 100 channels at each edge of the original spectra were excluded. Fourth- to tenth-order baselines were subtracted from each spectrum with 2000 to 3000 channels to avoid lines being truncated at the edge of subspectra, and then these subspectra were stitched together to resconstruct the complete spectrum. Since some of those subspectra also suffer contamination from time variable radio frequency interference (RFI), the channels showing RFI signals have been “flagged”, i.e. discarded from further analysis, resulting in some limited gaps in the averaged spectra. Furthermore, there were a few individual bad channels, which have also been eliminated.

2.2. Archival data

Here, we also present archival ALMA line survey data covering a frequency range of 214–247 GHz to study the distribution of molecules detected by our 1.3 cm line survey. The observations were carried out with 16 antennas on 2012 January 20, during ALMA’s science verification (SV) phase. Callisto and the quasar J0607-085 were used as the flux and phase calibrators, respectively. The spectral resolution is 0.488 MHz, which corresponds to ~0.6 km s^-1 at 230 GHz. Its primary beam size (~30′′) is slightly smaller than the beam size of the Effelsberg-100m in the λ ~ 1.3 cm band. The CLEAN algorithm in the Common Astronomy Software Applications package (CASA) was used to deconvolve the images by applying natural weighting. The synthesized beam size is 1.9′′× 1.4′′. In this work, we chose to use 11 unpublished transitions (see Table C.1) since several other lines and dust continuum emission from the ALMA data have been already studied (e.g., Neill et al. 2013b,a; Crockett et al. 2014; Wu et al. 2014; Hirota et al. 2015). The data are available to the public and can be accessed from the ALMA-SV website⁴.

Since the ALMA images without short spacings are only sensitive to small-scale structure, we also made use of the VLA C band continuum data⁵ (project AY191A) and the SCUBA−850 μm dust emission (Di Francesco et al. 2008) to trace the large-scale structure of Orion KL. In addition, we also compare the ALMA data with other studies to address the missing flux problem (see Appendix C).

3. Components of Orion KL

Figure 1a shows the large-scale structure of Orion KL. In Fig. 1b, we present an ALMA 230 GHz continuum image with the prominent sources near to or within Orion KL indicated. The positions of a few objects – Source n (a protostar coincident with a peculiar double 8.4 GHz radio continuum source), Source I (an embedded massive protostar without submillimeter and infrared continuum, but only SiO maser emission), SMA1 (a compact submillimeter source not detected in the infrared or cm regime) and the BN object (a massive protostar with a circumstellar disk), are based on previous studies (e.g., Menten & Reid 1995; Jiang et al. 2005; Beuther & Nissen 2008; Bally et al. 2011; Favre et al. 2011). The positions of IRc1, IRc2, IRc3, IRc4, IRc5, IRc6, and IRc7 infrared emission peaks are taken from Shuping et al. (2004). The millimeter continuum sources MM4, MM5, and MM6 are taken from Wu et al. (2014). In the λ ~ 1.3 cm band, the spectral features of Orion KL can stem from eight distinct components within our ~40′′ beam. These components are:

(1)
ionized gas from the foreground M42HII region (see Fig. 1a). Thecontribution can be traced by hydrogen and helium RRLs, whichhave typical velocitiesυ_lsr ~− 5 km s^-1 and large line widths Δυ> 15 km s^-1 (e.g., Wilson et al. 1997; Peimbert et al. 1992);
(2)
the photon-dominated region (PDR). It is the interface region between the molecular cloud and the foreground M42 HII region. It can be traced by carbon RRLs (υ_lsr ~ 9 − 10 km s^-1 and Δυ ~ 3 − 5 km s^-1) (e.g., Natta et al. 1994), which is because carbon ions can exist outside HII regions owing to their lower ionization energy (11.3 eV) with respect to that of hydrogen (13.6 eV) (e.g., Wyrowski et al. 1997; Wilson et al. 2009);
(3)
the “hot core” (HC) component. The HC is a well-known source, ~2′′ from IRc2 (see Fig. 1b). It has a high kinetic temperature and high density (T_kin> 150 K and n> 10⁷ cm^-3, e.g., Genzel & Stutzki 1989; Mangum & Wootten 1993; Schilke et al. 1997; Wilson et al. 2000; Wang et al. 2010; Goddi et al. 2011; Crockett et al. 2014). An extremely high kinetic temperature is also indicated by detections of high J metastable NH₃ transitions with energy levels up to 1900 K and a high rotational temperature of ~400 K (Wilson et al. 1993). The HC is believed to be heated by an explosive event (Bally et al. 2011; Zapata et al. 2011; Nissen et al. 2012). Its line profiles have velocities of υ_lsr ~ 3 − 6 km s^-1 and widths of Δυ ~ 5 − 12 km s^-1;
(4)
the “compact ridge” (CR) component. It is located ~12′′ southwest of the well-known HC (see Fig. 1b). It has densities of n ~ 10⁶ cm^-3 and temperatures of T_kin ~ 80 − 280 K (e.g., Mangum & Wootten 1993; Beuther et al. 2005; Persson et al. 2007; Wang et al. 2010; Favre et al. 2011). Studies of methanol and methyl formate emission suggest that the CR is externally heated (Wang et al. 2011; Favre et al. 2011). Its line profiles have velocities of υ_lsr ~ 7 − 9 km s^-1 and widths of Δυ ~ 3 − 6 km s^-1;
(5)
the “hot core south” (HC(S)) component. The component was first noticed by studies of HDO (Neill et al. 2013b; Crockett et al. 2014) and is likely to originate 1′′ south of the well-known HC submillimeter continuum peak (see Fig. 1b). It shows a spectral feature with υ_lsr ~ 6.5 − 8 km s^-1 and Δυ ~ 5 − 10 km s^-1;
(6)
the “extended ridge” (ER) component. It represents the ambient gas of Orion KL. It has densities of n ~ 10⁴ − 10⁶ cm^-3 and temperatures of T_kin< 60 K showing spectral features of υ_lsr ~ 8 − 10 km s^-1 and Δυ ~ 2 − 5 km s^-1 (e.g., Blake et al. 1986; Schilke et al. 1997; Persson et al. 2007);
(7)
the “plateau” component. The component is characterized by velocities υ_lsr ~ 6 − 12 km s^-1 and large line widths Δυ> 20 km s^-1, which are attributed to outflows. Orion KL is known to contain at least two outflows, which are historically referred to as the low velocity flow (LVF) and the high velocity flow (HVF; e.g., Genzel et al. 1981). The LVF is oriented along the SW-NE direction and is thought to be driven by radio source I (e.g., Genzel & Stutzki 1989; Menten & Reid 1995; Greenhill et al. 1998; Plambeck et al. 2009). It has spectral features of υ_lsr ~ 5 km s^-1 and Δυ ~ 18 km s^-1. More extended than the LVF, the HVF is oriented along the SE-NW direction and is thought to be driven by the submillimeter source SMA1 or the dynamical decay of a multistar system involving radio source I, source n and the BN object that caused the explosive event described in (3) above (e.g., Beuther & Nissen 2008; Bally et al. 2011; Goicoechea et al. 2015). It has spectral features of υ_lsr ~ 10 km s^-1 , and the line widths can be very large and reach velocities up to 150 km s^-1;
(8)
the millimeter continuum sources MM4, MM5, and MM6 taken from Wu et al. (2014). Based on their ¹³CH₃CN studies, MM4, MM5, and MM6 are found to have rotational temperatures of 182 $\begin{matrix} + 81 \\ -43 \end{matrix}$ $\hbox{$^{+81}_{-43}$}$ K, 157 $\begin{matrix} + 43 \\ -28 \end{matrix}$ $\hbox{$^{+43}_{-28}$}$ K, and 119 $\begin{matrix} + 18 \\ -14 \end{matrix}$ $\hbox{$^{+18}_{-14}$}$ K, respectively. The ¹³CH₃CN spectral features are: MM4 with velocities of υ_lsr ~ 5.5 km s^-1 and widths of Δυ ~ 5.9 km s^-1; MM5 with velocities of υ_lsr ~ 10.1 km s^-1 and widths of Δυ ~ 5.4 km s^-1; MM6 with velocities of υ_lsr ~ 7.5 km s^-1 and widths of Δυ ~ 5.7 km s^-1. From Fig. 1b, MM4 is close to IRc7 while MM6 is close to IRc6.

4. Results

4.1. Line identifications

The line identification was performed with the help of the JPL⁶, CDMS⁷, and splatalogue⁸ databases as well as with the online Lovas line list⁹ for astronomical spectroscopy (Müller et al. 2005; Pickett et al. 1998; Lovas & Dragoset 2004). However, the rest frequencies of RRLs are not contained in these databases, and are thus calculated with the Rydberg formula $ν_{n_{2} \to n_{1}} = R (\frac{1}{n_{1}^{2}} - \frac{1}{n_{2}^{2}}), n_{1} < n_{2},$ $\begin{equation} \label{redberg} \nu_{n_{2}\to n_{1}} = R\left(\frac{1}{n_{1}^{2}}-\frac{1}{n_{2}^{2}}\right) ,~n_{1} < n_{2} , \end{equation}$ (1)where Δn = n₂ − n₁. n₂ and n₁ are the principle quantum numbers of the upper and lower states, and R is the Rydberg constant that is equal to 3.28805129 × 10¹⁵ Hz for hydrogen and 3.28939118 × 10¹⁵ Hz for helium (Wilson et al. 2009). The nomenclature for RRLs is based on n₁ and Δn. For example, H71α corresponds to a hydrogen RRL with n₁ = 71, n₂ = 72, and Δn = 1. The Greek alphabet corresponds to Δn, e.g., α corresponds to Δn = 1, β corresponds to Δn = 2, etc.

Fig. 2

Overview of the 1.3 cm line survey toward Orion KL with strong lines marked. The displayed frequency scale is based on the Local Standard of Rest velocity 0 km s^-1.

A line is identified as real if it exhibits a 3σ feature in more than three adjacent raw channels. On the other hand, there are a few narrow lines that are also identified as real even though their intensity is above 3σ for fewer than three adjacent raw channels. Such lines show up as doublets (e.g., CH₃OCHO, CH₃OCH₃), which confirms our assignment. Identified lines are listed in Table 2. Given that line profiles are complex in Orion KL, different methods are employed to fit our observed lines. For RRLs, we simply use one single Gaussian component to derive their observed properties. For molecular lines showing nuclear quadrupole hyperfine structure (hfs, like HC₃N, CH₃CN, HNCO, NH₃, etc.), we used the HFS fitting routine embedded in CLASS to derive the observed properties (including the optical depth) while least-square fits of Gaussians are performed for the other molecular transitions. On the other hand, the observed lines of sulfur-bearing molecules are composites of different distinct components, so here the approach we take is to adopt a range of velocities and line widths of different components to separate their contributions. This is based on general knowledge of the source (see Sect. 3). The NH₃ emission also originates in several components, but here the situation is even more complex. In this case, multiple component fitting becomes unrealistic. Therefore, we only give the integrated intensities for those NH₃ transitions that have a heterogeneous origin. Although Cα lines are blended with Heα lines, we present double Gaussian fits to separate their relative contributions to obtain the observed properties of Cα lines since little information about Cα lines exists in the literature. For other blended transitions, we do not separate the lines but simply integrate the whole line profiles to obtain upper limits to the integrated intensities of individual spatial subcomponents.

The observed line parameters for different species are displayed in Tables A.1−A.6, where the rms noise σ given is for 61 kHz wide channels. Figure 2 shows an overview of the 1.3 cm spectral line survey toward Orion KL. In the survey, we find 261 lines, yielding an average line density of about 32 spectral features per GHz above 3σ. The line density is much lower than that at millimeter wavelengths (e.g., Tercero et al. 2010). All lines are identified. The identified lines include 164 RRLs and 97 molecular lines. The RRLs, from hydrogen, helium, and carbon, originate in the ionized material of the Orion Nebula, part of which is covered by our beam. The molecular lines are assigned to 13 different molecules including rare isotopologues. Twenty-three lines from already known species are detected for the first time in the interstellar medium (marked by an “N” in the last column of Table 2), and 19 molecular lines or carbon RRLs are blended with neighboring hydrogen or helium RRLs. From Fig. 2, we can see that the spectrum is dominated by strong lines arising from four species: the super strong H₂O (6_1,6 − 5_2,3) maser with S_ν> 30 000 Jy, the series of CH₃OH maser lines starting at ≈25 GHz with S_ν>10 Jy, low J metastable transitions of NH₃ with S_ν> 5 Jy and Hα lines with S_ν> 2 Jy. Figure B.1 shows the observed spectrum in more detail. Each panel covers ~500 MHz with a 10 MHz overlap between adjacent panels so that lines truncated in one panel will not be truncated in the neighboring panel. In addition, Figs. B.2–B.16 shows zoomed-in plots of all identified lines.

4.2. Radio recombination lines

From Fig. B.1 and Table 2, we find that the frequency range is mainly occupied by RRLs. Among the 164 RRLs, there are 116 hydrogen RRLs, 39 helium RRLs, and 9 carbon RRLs. Their spectra are shown in Figs. B.2−B.10, and their observed line parameters are given in Table A.1. Fifty-six hydrogen RRLs with Δn ≥ 6 are detected. Furthermore, hydrogen RRLs with Δn as high as 11 are detected in this survey, and hydrogen RRLs with even higher Δn (≤ 25) are reported by Bell et al. (2011). Carbon RRLs are all blended with helium RRLs since line widths of RRLs are broad and the separations between the rest frequencies of carbon RRLs and helium RRLs with the same n₁ and Δn are small. Furthermore, the velocity differences shift the carbon RRLs slightly into the He RRLs (See Sect. 3 for the differences with respect to radial velocities.).

From Fig. 3a, we can see that line widths (~24 km s^-1) of hydrogen RRLs are broader than those (~16 km s^-1) of helium RRLs. The line broadening of RRLs is thought to arise from thermal motions and turbulence (Wilson et al. 2009). The observed line width, Δυ, has both a thermal and a turbulent contribution, $Δ υ = \sqrt{Δ υ_{th}^{2} + Δ υ_{tur}^{2}}$ $\hbox{$\Delta \upsilon = \sqrt{\Delta \upsilon_{\rm th}^{2}+\Delta \upsilon_{\rm tur}^{2}}$}$ , where Δυ_th is the thermal width, and Δυ_tur is the line width caused by turbulent flows. Taking the electron temperature of the foreground HII region as 8300 K (Wilson et al. 1997) and following a Maxwell-Boltzmann velocity distribution, the thermal widths of hydrogen and helium are estimated to be 19.5 km s^-1 and 9.8 km s^-1with the formula $Δ υ_{th} = 2 \sqrt{2 \ln 2} (\frac{kT}{m})^{1 / 2}$ $\hbox{$\Delta \upsilon_{\rm th} = 2\sqrt{2{\rm ln}2}(\frac{{k}T}{m})^{1/2}$}$ , where k is the Boltzmann constant, T the electron temperature, and m the mass of a particle. Consequently, the average hydrogen and helium turbulent line widths are estimated to be 14.1 ± 1.0 km s^-1 and 12.4 ± 1.4 km s^-1, respectively. From Fig. 3b, we can see that most points lie below the red dashed line where the turbulent line widths are equal for hydrogen and helium. This suggests that hydrogen turbulent line widths might be slightly larger than those of helium. However, the large errors make it uncertain. If the difference can be confirmed, that may indicate that hydrogen RRLs trace a larger portion of ionized gas than helium RRLs in Orion KL. The exciting stars of the HII region have T_eff ~ 4 × 10⁴ K (Baldwin et al. 1991). This is higher than the critical T_eff of 37 000 K, below which helium is no longer appreciably ionized (Wilson et al. 2009). The innermost parts of their HII regions should therefore show ionized HeII. The outer parts, however, are devoid of HeII, while still containing ionized hydrogen. The larger volume leads to a larger line width. Nevertheless, the bulk of the discrepancy between observed line widths of hydrogen and helium RRLs stems from their thermal contributions.

Fig. 3

a) Relationship between the observed line widths of hydrogen (abscissa) and helium (ordinate) RRLs. The red dashed lines represent the unweighted mean values of observed line widths of hydrogen and helium RRLs, respectively. b) The relationship between the turbulent line widths derived from hydrogen and helium RRLs. The red dashed line connects points where turbulent line widths for hydrogen and helium are equal.

Assuming local thermodynamic equilibrium (LTE) and following the formulae (6.24)–(6.27) of Brocklehurst & Seaton (1972), the intensity ratios of RRL pairs with different Δn from the same atom at neighboring frequencies (e.g., H88β/H70α) are found to be $\frac{T_{Δ n_{1}}}{T_{Δ n_{2}}} = \frac{Δ n_{1} K (Δ n_{1})}{Δ n_{2} K (Δ n_{2})},$ $\begin{equation} \frac{T_{\Delta n_{1}}}{T_{\Delta n_{2}}} = \frac{\Delta n_{1}K(\Delta n_{1})}{\Delta n_{2}K(\Delta n_{2})}, \end{equation}$ (2)where Δn₁ and Δn₂ represent different Δn defined in formula (1), and ΔnK(Δn) is given in Table I of Brocklehurst & Seaton (1972). From Fig. 4, we find that observed ratios from Hβ/Hα, Hγ/Hα, and Hδ/Hα are consistent with LTE ratios. Based on a table of the departure coefficients b_n provided by Salem & Brocklehurst (1979), we find that all b_n factors are larger than 0.94 and smaller than unity, when adopting an electron temperature of 8300 K and an electron density of 1 × 10⁴ cm^-3 (Wilson et al. 1997). Therefore, we suggest that the LTE deviations are very small for these RRLs in Orion KL.

4.3. Molecular lines

In this section, we focus on the origin of detected molecular lines according to the line parameters obtained by Gaussian fits or hfs fits in our survey, as well as by spatial distributions of their corresponding transitions at higher frequencies obtained from the ALMA-SV data (see Table C.1 and Fig. C.1).

Fig. 4

Comparison of observed and LTE ratios of recombination lines. The observed ratios for hydrogen and helium RRLs are indicated with circles and pentagrams. The red dashed lines represent the LTE ratios, while the blue dashed lines represent the LTE ratios with the departure coefficients b_n corrected. Each ratio is given below the respective abscissa.

4.3.1. NH₃, ¹⁵NH₃, and NH₂D

In our band, we find 7 metastable NH₃ transitions, 27 non-metastable NH₃ transitions, 7 metastable ¹⁵NH₃ transitions, 3 non-metastable ¹⁵NH₃ transitions, and 2 NH₂D transitions (see Table A.2 and Fig. B.11). It is worth noting that the non-metastable ¹⁵NH₃ transitions are detected in Orion KL for the first time. The line widths of non-metastable ¹⁵NH₃ transitions become narrower with increasing J, while there is no such trend for the metastable ¹⁵NH₃ transitions (see Fig. B.12 and Table A.2). This suggests that higher-J non-metastable transitions may trace denser regions where turbulent motions become less dominant. We note that the errors of the line widths (~1 km s^-1) are large compared to the differences between these line widths, so further observations are needed to confirm the trend. Based on previous VLA ammonia observations (e.g., Genzel et al. 1982; Goddi et al. 2011), NH₃ emission is mainly found in the HC, the Plateau, MM4, MM5, and MM6, but barely in the CR. Velocities around 6 km s^-1 and large line widths suggest that ¹⁵NH₃ mainly comes from the HC. One of the two NH₂D transitions belongs to the para, the other to the ortho species ¹⁰, which have already been reported by Walmsley et al. (1987). The intensity of NH₂D (4_1,4a − 4_0,4s) from our observations is lower than theirs, but is still within their uncertainties. NH₂D (3_1,3s − 3_0,3a) is blended with NH₃ (8,5) at 18808.5 MHz and H118ϵ at 18808.4 MHz (see Fig. B.12). From Fig. C.1, NH₂D emission is found in the HC, the HC(S), and in MM4. NH₂D (3_1,3s − 3_0,3a) and (4_1,4a − 4_0,4s) are considered to arise from the HC in the following analysis.

4.3.2. H₂O and HDO

We find one H₂O line and one HDO line (see Table A.6). The H₂O (6_1,6 − 5_2,3) maser is blended with NH₃ (3, 1), and shows multiple velocity components that have peak intensities of >1000 Jy (see Fig. B.16). A previous deuterated water study has shown that HDO mainly comes from the HC(S) (see Fig. 6 of Neill et al. 2013b). Although HDO (3_2,1 − 4_1,4) is blended, it clearly peaks at ~7 km s^-1, which coincides with the velocity of the HC(S). This also indicates that the deuterated water arises mainly from the HC(S).

4.3.3. Sulfur-bearing molecules, SO₂ and OCS

We find five transitions of SO₂ and one transition of OCS in the survey. Most of the unblended lines can be fit with two velocity components (see Fig. B.14 and Table A.4). From Fig. C.1, we can see that the SO₂ emission around 6 km s^-1 mainly comes from the HC. Thus, the two velocity components are assigned to the HC and the plateau. But for SO₂ (8_1,7–7_2,6), it cannot be fitted with two Gaussian components. There is an additional narrow component that has a velocity of ~8 km s^-1 and a line width of ~2 km s^-1 when we fit the line with three Gaussian components. It may result from emission from the CR, an outstanding emitter of transitions at low upper level energy, but it is narrower than the typical SO₂ line width (3–6 km s^-1) of the CR (Esplugues et al. 2013b). Alternatively, it may be affected by population inversion since the narrow component is not detected in other SO₂ transitions in our band and the 1.3, 2 mm, and 3 mm bands (Esplugues et al. 2013b). Furthermore, SO₂ is extremely enhanced by shocks so that it becomes the strongest coolant in the 607–725 GHz survey which outperforms CO by a large factor (Schilke et al. 2001).

OCS (2–1) also shows two components, but the fitted velocity and line width are different from those derived from SO₂ (see Table A.4). The line width of the plateau component from OCS is much narrower than found for SO₂. This is probably because OCS (2–1) is only excited in the LVF. Based on the distributions of OCS (see Fig. C.1), the other component of OCS (2–1) is assigned to the HC(S) according to the fitted velocity and line width.

4.3.4. Cyanopolyynes, HC₃N and HC₅N

Cyanopolyynes are well studied and abundant in many different astronomical environments, such as the late-type carbon star IRC +10216 and the cold starless core TMC-1 (e.g., Cernicharo et al. 2000; Kaifu et al. 2004), whereas they are not prominent in Orion KL (Esplugues et al. 2013a). BIMA observations show that toward Orion KL, HC₃N emission originates in the HC (Wright et al. 1996), but the origin of HC₅N is still a puzzle. In our survey, we find one HC₃N line and two HC₅N lines (see Fig. B.16). Based on the measured velocities and line widths (see Table A.6), HC₅N seems to come from the CR or the HC(S) or the ER. From Fig. C.1, we confirm that HC₃N is mainly from HC, while HC₅N is not detected in the archival ALMA data. Nevertheless, it would be very surprising if HC₃N and HC₅N did not share the same spatial origin in Orion KL. Assuming that HC₅N originates in the same region as HC₃N, the HC₅N/HC₃N abundance ratio is estimated to be (2.8 ± 0.7) × 10^-2 according to Table 3. This is slightly lower than the (7 ± 4) × 10^-2 of Esplugues et al. (2013a). HC₅N (7 − 6) is inside the frequency range covered by our survey. With the rotational temperatures and column densities listed in Table 3, we can estimate the peak intensity of HC₅N (7 − 6) by assuming a low opacity, a Gaussian profile, a line width of 5 km s^-1, and a source size of 10′′. The peak intensity is estimated to be about 0.012 Jy, so its non-detection can be expected. On the other hand, we do not detect HC₇N and HC₉N transitions in this band which are detected in large numbers toward TMC-1 and IRC +10216 (Kaifu et al. 2004; Gong et al. 2015). HCN υ₂ = 1 direct l-type transitions detected originally in the protoplanetary nebula CRL 618 (Thorwirth et al. 2003) and selected high-mass star-forming regions (Thorwirth 2001), as well as the prototypical starburst galaxy Arp 220 (Salter et al. 2008), are also not detected in this survey. A tentative detection of the J = 9 direct l-type transition toward Orion KL reported earlier (Thorwirth 2001) could not be confirmed here.

Table 3

Column densities and rotational temperatures of the detected molecules compared with previous studies.

4.3.5. CH₃OH and ¹³CH₃OH

In this survey, we find seventeen CH₃OH (v_t = 0) transitions, four torsionally excited CH₃OH (v_t = 1) transitions, and one ${}^{13}{CH}_{3} OH$ $\hbox{$^{13}{\rm CH}_{3}{\rm OH}$}$ (v_t = 0) transition. The J₂ − J₁E methanol masers with J = 2 ... 10, already reported (e.g., Menten et al. 1988a), are all detected. From Fig. 5, our measurements show that the peak intensities of these masers increase with principle quantum number J from 2 to 6 and decrease with J from 6 to 10. This is consistent with the result of Menten et al. (1986), although the J₂ − J₁E methanol masers with J = 8, 9 are not included in their analysis. The newly detected CH₃OH (26₂ − 26₁E) line, belonging to the J₂ − J₁E class, might not to be affected by population inversion, since it agrees well with the fit to the rotational diagram to CH₃OH transitions (see Sect. 5.1). From Fig. C.1, we can see that methanol exists in the HC, the HC(S), the CR, MM4, MM5, and MM6. Based on the fitted line widths and velocities (see Table A.3), we suggest that the transitions detected by our survey mainly arise from the CR.

Fig. 5

Observed peak intensities of CH₃OH (J₂ − J₁E) masers as a function of rotational quantum number J.

4.3.6. HNCO

From Fig. C.1, HNCO is found to exist mainly in the HC, the HC(S), and MM4 while only weak emission arises from MM5, MM6, and the CR. The HNCO emission in HC(S) stands out around velocities ranging from 6–10 km s^-1. In combination with the fitted parameters of HNCO (1–0) (see Table A.6), we suggest that HNCO (1–0) mainly originates in the HC(S). HNCO is believed to be enhanced in shocked regions (Zinchenko et al. 2000), which may indicate that the HC(S) component is formed via shocks due to the interaction between the outflowing gas and ambient clouds.

4.3.7. H₂CO

In this survey, we find one transition of H₂CO, the 9_2,7 − 9_2,8 line, which is from the para state of H₂CO. From Fig. C.1, H₂CO is found to exist in the HC, the HC(S), the CR, the Plateau, MM5, and MM6. Based on the fitted parameters (see Table A.6), H₂CO (9_2,7 − 9_2,8) may arise from the HC(S) or the CR. On the other hand, H₂CO (9_2,7 − 9_2,8) has an upper energy of 205 K, so we suggest that H₂CO (9_2,7 − 9_2,8) is mainly from the HC(S).

4.3.8. CH₃CN and CH₃CH₂CN

One CH₃CN transition and one CH₃CH₂CN transition are detected in our survey. Based on Fig. C.1, we find that the CH₃CN emission arises from many spatial components, among which the HC is the strongest, while CH₃CH₂CN emission is mainly from the HC, followed by MM4. Inspecting the fitted parameters for CH₃CN (1₀ − 0₀) and CH₃CH₂CN (3_1,3 − 2_1,2) (see Table A.6), we suggest that the two transitions are also mainly from the HC. The line profile of CH₃CH₂CN (3_1,3 − 2_1,2) is broad, because its hfs lines are blended. Although many other CH₃CH₂CN transitions with similar upper level energies fall in this band, they are expected to be much weaker since their intrinsic strengths are much lower.

4.3.9. CH₃OCHO and CH₃OCH₃

In our survey, we find twelve CH₃OCHO transitions and two CH₃OCH₃ transitions (see Tables A.5 and A.6). Based on the fitted velocities and narrow line widths, we suggest that CH₃OCHO and CH₃OCH₃ originate in the CR. From Fig. C.1, we find that the emissions of CH₃OCHO and CH₃OCH₃ mainly come from the CR, followed by MM5 and MM6.

4.3.10. A brief summary

The origin of molecules detected by our 1.3 cm line survey is briefly summarized in Table 4, demonstrating that chemical differentiation exists in Orion KL. This differentiation is roughly consistent with the fact that nitrogen-bearing molecules are favored in the HC while oxygen-bearing molecules are more predominant in the CR (e.g., Beuther et al. 2005; Friedel & Snyder 2008).

5. Discussion

5.1. Rotational temperatures, column densities and fractional abundances relative to H₂

Given that the density in Orion KL is relatively high (see Sect. 3), LTE should be a good approximation. In the 1.3 cm wavelength range, the continuum emission is dominated by free-free emission. Toward Orion KL, free-free emission originates mainly in its foreground HII region M42 and is optically thin in the 18–26 GHz band (Mezger & Henderson 1967; Terzian & Parrish 1970), so the free-free continuum emission of Orion KL itself is simply assumed to be zero in the radiative transfer equation to estimate the opacity of molecular transitions at 1.3 cm wavelength. Thus, the opacity at the line center can be estimated via a formula given by Wilson et al. (2009), $τ = - \ln (1 - \frac{T_{mb}}{f (J (T_{rot}) - J (T_{bg}))}),$ $\begin{equation} \label{tau} \tau = - {\rm ln} \left(1- \frac{T_{\rm mb}}{f(J(T_{\rm rot})-J(T_{\rm bg}))} \right) , \end{equation}$ (3)where T_mb is the observed peak intensity, T_rot the rotational temperature, and T_bg the 2.73 K cosmic microwave background brightness temperature. Here, f is the filling factor $f = θ_{s}^{2} / (θ_{s}^{2} + θ_{beam}^{2})$ $\hbox{$f=\theta^{2}_{\rm s}/(\theta^{2}_{\rm s}+\theta^{2}_{\rm beam})$}$ , where we assume a Gaussian source and beam width FWHM of θ_s and θ_beam, respectively, and $J (T) = \frac{hν}{k} \frac{1}{e^{hν / kT} - 1}$ $\hbox{$J(T)=\frac{h\nu}{k} \frac{1}{{\rm e}^{h\nu/kT}-1}$}$ , h is the Planck constant, ν the rest frequency of the analyzed transition, and k the Boltzmann constant. If we take a filling factor of 0.06 (corresponding to an assumed source size of 10′′) and a rotational temperature of 100 K, a peak intensity of 1 K (corresponding to ~0.7 Jy at 22 GHz) is equivalent to an opacity of 0.18 at the line center. When a higher rotational temperature is taken, the opacity becomes even lower. Thus, most of the detected thermally excited molecular lines appear to be optically thin with the exception of the low J NH₃ transitions. On the other hand, lines with hfs allow for determinations of opacity as well. Based on the HFS fitting routine in CLASS, we confirm that HC₃N (2−1), CH₃CN (1₀ − 0₀), HNCO (1_0,1 − 0_0,0), and CH₃CH₂CN (3_1,3 − 2_1,2) are optically thin (see Table A.6).

Below, we use rotational diagrams to roughly estimate rotational temperatures and column densities. The standard formula used here is $\ln (\frac{3 kW}{8 π^{3} ν μ^{2} S}) = \ln (\frac{N_{tot}}{Q (T_{rot})}) - \frac{E_{u}}{k T_{rot}},$ $\begin{equation} \label{f1} {\rm ln}\left(\frac{3kW}{8\pi^{3}\nu\mu^{2}S}\right)~=~{\rm ln}\left(\frac{N_{\rm tot}}{Q(T_{\rm rot})}\right)-\frac{E_{\rm u}}{kT_{\rm rot}}, \end{equation}$ (4)where k is the Boltzmann constant, W the integrated intensity, ν the rest frequency, μ the permanent dipole moment, S the transition’s intrinsic strength, N_tot the total column density, T_rot the rotational temperature, Q the partition function, and E_u is the upper level energy of the transition. The values of Q and μ are taken from the CDMS and JPL catalogs. The partition function for CH₃OH and ¹³CH₃OH has included the vibrational state v_t = 0 and v_t = 1 (Dr. Christian P. Endres, priv. comm.). For molecules with at least two transitions, we use least-square fits to the rotational diagrams to derive their rotational temperatures and column densities. For molecules with only one transition detected, we fix rotational temperatures, according to the values from the literature or chemically related molecules, to derive their column densities (see Table 3). But for HC₅N, the fitted rotational temperature is about 1 K, which is not reasonable. This is because we only detect two HC₅N transitions, and their upper level energy difference (~1 K) is very small. Thus, rotational temperatures of 30 K and 200 K are taken to derive the column density of HC₅N, assuming that HC₅N comes from the CR and HC, respectively.

Following previous studies (Neill et al. 2013a; Crockett et al. 2014), we assume that all molecules with the same origin have the same source size to simplify the calculation. Adopted source sizes for each component are also based on Table 2 of Crockett et al. (2014), and the parameters of the HC(S) are assumed to be the same as those of the HC. With the adopted source sizes, the observed intensities are corrected for beam dilution by dividing by the filling factor to derive source-averaged column densities. The H₂ column densities adopted in Crockett et al. (2014) are used to calculate fractional abundances of the detected molecules.

Figure 6 shows rotational diagrams for molecules with more than two transitions detected in our survey. For NH₃, we only fit the non-metastable transitions since metastable transitions consist of several components, and decompositions based on the velocity information may be less accurate. In the fitting process, we separate non-metastable transitions into four classes that are labeled as J = K + 1, J = K + 2, J = K + 3, and J = K + 4. For J = K + 1, NH₃ (2, 1), (3, 2), (4, 3), (5, 4), (6, 5), and (7, 6) are ignored in the fitting process since they are potentially contaminated by emission from other components rather than only from the HC component (see Fig. B.11). From the fitted result (see Fig. 6) where the ortho-to-para ratio is set to unity (based on the analysis of ¹⁵NH₃, see the discussion below), we find that only NH₃ (2, 1) and (3, 2) deviate significantly, while all others lie along the fitted line which suggests that other J = K + 1 transitions are also dominated by the HC component. Nevertheless, if NH₃ (8, 7), (9, 8), (10, 9), and (11, 10) are not optically thin, the derived rotational temperature will be overestimated, and the derived column density will be underestimated.

Fig. 6

Rotational diagrams for NH₃, ¹⁵NH₃, CH₃OCHO, SO₂, and CH₃OH. The circles and squares are explained by the legends in the upper right of each panel. A calibration error of 20% has been included in the error bars. In the NH₃ panel, least-square fits to J + 1, J + 2, J + 3, and J + 4 are indicated by a red dashed line, a red dotted line, a red dot-dashed line, and a red solid line. In the ¹⁵NH₃ panel, the para and ortho transitions of ¹⁵NH₃ are black and blue; the least-square fits to para and ortho metastable ¹⁵NH₃ transitions and non-metastable ¹⁵NH₃ transitions are indicated by a red dashed line, a red dotted line and a red dot-dashed line, respectively. In the CH₃OH and CH₃OCHO panels, the dashed and dotted lines represent fits to A type and E type transitions, respectively. In the SO₂ panel, dashed and dotted lines represent fits to its plateau and its HC component.

We also find that the fits for J = K + 2 and J = K + 3 have rotational temperatures similar to J = K + 1. However, the fit for J = K + 4 has a much lower rotational temperature. Furthermore, the derived NH₃ column densities go up from J = K + 1 to J = K + 4. This suggests that the four classes may trace different NH₃ emission regions. However, it is surprising that the J = K + 4 class has the highest column density among the four since the J = K + 4 transitions require particularly high excitation. We also note that the fit for the J = K + 4 class may have large uncertainties since we only have three data points.

To determine the ortho-to-para ratio of ammonia, ¹⁵NH₃ is more suitable since the optical depths are much smaller than those of NH₃. Assuming that this rare isotopologue was formed in the same region as its main species, we fit its para and ortho states of metastable transitions separately (see Fig. 6). The ortho-to-para ratio is estimated to be 0.99 ± 0.34, in accord with the result ( $0. 7_{-0.3}^{+ 0.5}$ $\hbox{$0.7^{+0.5}_{-0.3}$}$ ) of Hermsen et al. (1985), which agrees very well with the fact that the ortho-to-para ratio approaches the statistical ratio of unity when ammonia is equilibrated under high kinetic temperature conditions (e.g., Takano et al. 2002). On the other hand, we find that metastable ¹⁵NH₃ has a higher excitation temperature than non-metastable ¹⁵NH₃ although only three points have been fitted for the non-metastable ¹⁵NH₃ (see Fig. 6 and Table 3). This is very likely because the metastable line excitation is dominated by collisions, while the excitation of the non-metastable transitions follows the radiation field. In the very dense medium, collisional (T_kin) and radiative (T_rad) temperatures could be the same. However, T_rad may suffer from beam dilution effects, which results in lower rotational temperatures. Alternatively, non-metastable ¹⁵NH₃ transitions trace higher densities than its metastable counterparts and thus represent the temperature of higher excited inner regions than metastable transitions. Meanwhile, the HC is believed to be externally heated by an explosive event (Bally et al. 2011; Zapata et al. 2011; Nissen et al. 2012), so the innermost region could be colder. Therefore, non-metastable transitions of ¹⁵NH₃ can be expected to have a lower rotational temperature than its metastable transitions. This is also consistent with the fact that the J = K + 4 transitions of ¹⁴NH₃, tracing denser regions, have a lower rotational temperature.

Fig. 7

He⁺/H⁺ abundance ratios derived from RRLs as a function of rest frequency. The abundance ratios derived from Heα/Hα, Heβ/Hβ, Heγ/Hγ, and Heδ/Hδ are marked with red filled circles, blue filled triangles, green filled squares and black filled diamonds, respectively. The dashed line represents the sigma-weighted mean value of He⁺/H⁺ abundance ratios.

No CH₃OH maser lines were included in the fit owing to their deviation from LTE. CH₃OH (2₁ − 3₀E) was also not included because this line comprises several velocity components due to its low upper level energy (E_u/k = 28 K). We fit the A type and the E type of CH₃OH and CH₃OCHO separately. Consequently, the CH₃OCHO A/E abundance ratio is estimated to be 1.38 ± 0.26, so slightly higher than the expected value of unity, while the CH₃OH A/E abundance ratio is estimated to be 0.77 ± 0.33, so slightly lower than the expected value of unity. Nevertheless, the differences could be caused by uncertainties. For SO₂, we fit its HC component and its plateau component separately. We find that the HC component has a higher rotational temperature and a higher column density than its plateau component.

The resulting rotational temperatures and column densities, together with results obtained with single-dish telescopes from the literature, are given in Table 3. We find fitted rotational temperatures ranging from 69 to 209 K and derived molecular column densities ranging from 9.6 × 10¹³ to 1.9 × 10¹⁸ cm^-2. This results in fractional abundances spanning more than four orders of magnitude. Taking beam dilution into account, we find that most of our results agree well with previous studies based on (sub)millimeter lines except for SO₂, HC₃N, HC₅N, CH₃CN, and HNCO. Our SO₂ column density for the HC is less than that of Crockett et al. (2014) by a factor of 6. This is because our fitted rotational temperature is much lower and the SO₂ lines of Crockett et al. (2014) at higher frequencies, trace higher temperatures. The derived column densities of HC₃N, HC₅N, CH₃CN, and HNCO are higher than those from previous studies listed in Table 3. The upper level energies of their transitions are very low, so they are probably contaminated by other components not being part of the HC. Thus, our assumed source size of 10″ could be an underestimate, which results in an overestimated column density. The derived column densities should then be upper limits for the average column density of those molecules. Alternatively, previously reported column densities were derived from (sub)millimeter lines which may suffer from opacity effects, leading to an underestimate of their column densities.

5.2. Elemental and isotopic abundance ratios in Orion KL

The He/H, ¹²C/¹³C, ¹⁴N/¹⁵N, and D/H abundance ratios are fundamental parameters for the study of cosmology and the evolution of galaxies (e.g., Boesgaard & Steigman 1985; Wilson & Rood 1994; Steigman 2007).

5.2.1. He/H

Benefiting from a large number of RRLs, we can make use of hydrogen and helium RRL pairs to estimate the He⁺/H⁺ abundance ratio that can be used to evaluate the He/H abundance ratio. Under LTE conditions, the He/H abundance ratio can be calculated via the integrated intensity of Gaussian fits to the hydrogen and helium RRLs (e.g., Churchwell et al. 1974): $y = \frac{N (He)}{N (H)} ≃ \frac{1}{R} \times \frac{\int_{Ω_{s}} dΩ \int S_{ν} (H e^{+}) d υ}{\int_{Ω_{s}} dΩ \int S_{ν} (H^{+}) d υ},$ $\begin{equation} \label{f2} y= \frac{N({\rm He})}{N({\rm H})} \simeq \frac{1}{R} \times\frac{\int _{\Omega_{\rm s}}{\rm d}\Omega \int S_{\nu}({\rm He^{+}}){\rm d}\upsilon}{\int_{\Omega_{\rm s}}{\rm d}\Omega \int S_{\nu}({\rm H^{+}}){\rm d}\upsilon} , \end{equation}$ (5)where R is the ratio of volumes of He⁺ and H⁺ Strömgren spheres, weighted by the square of the proton density, Ω_s is the source solid angle, ${}^{\int}S_{ν} ({He}^{+}) d υ$ $\hbox{$\int S_{\nu}({\rm He}^{+}){\rm d}\upsilon$}$ and ${}^{\int}S_{ν} (H^{+}) d υ$ $\hbox{$\int S_{\nu}({\rm H}^{+}){\rm d}\upsilon$}$ are the observed integrated intensities of corresponding helium and hydrogen RRLs. Assuming that the HII region has the same size as the HeII region, formula (5) becomes $y = \frac{N (He)}{N (H)} ≃ \frac{\int S_{ν} (H e^{+}) d υ}{\int S_{ν} (H^{+}) d υ} \cdot$ $\begin{equation} \label{f3} y= \frac{N({\rm He})}{N({\rm H})}\simeq \frac{\int S_{\nu}({\rm He^{+}}){\rm d}\upsilon}{\int S_{\nu}({\rm H^{+}}){\rm d}\upsilon} \cdot \end{equation}$ (6)In this work, there are 28 unblended pairs that are from Hα/Heα, Hβ/Heβ, Hγ/Heγ, and Hδ/Heδ. The rest frequencies of these pairs are only about 7 MHz apart, so their ratios should be free of uncertainties related to pointing accuracy, calibration errors, and different beam widths. Figure 7 shows the derived He/H abundance ratios according to formula (6). We find that uncertainties increase from Hα/Heα to Hδ/Heδ. Thus, we take the sigma-weighted mean value as the He/H abundance ratio that is estimated to be (8.7 ± 0.7)%. This agrees well with previous studies (Peimbert et al. 1992; Wilson & Rood 1994; Quireza et al. 2006) and is slightly lower than the (9.1 ± 0.5)% from previous H66α and He66α measurements, which were pointed at the HII maximum (Thum et al. 1980). Therefore, we confirm that the He/H abundance ratio is comparable to the primordial He/H abundance ratio of ~8.3% due to Big Bang nucleosynthesis (BBN; Olive & Skillman 2004; Steigman 2007) and the solar value (~9.8%; Wilson & Rood 1994), which indicates that the He/H abundance ratio is mainly determined by BBN but little affected by stellar nucleosynthesis. Furthermore, the dispersion of the ratios derived from Hα/Heα pairs to Hδ/Heδ pairs is small (see Fig. 7), suggesting that the departure coefficients, the b_n factors, are nearly identical for hydrogen and helium. This agrees well with theoretical predictions (e.g., Storey & Hummer 1995).

5.2.2. ¹²C/¹³C

We use the total column density of CH₃OH and ¹³CH₃OH, including both the A and E type, to estimate the ¹²C/¹³C ratio in the CR. For CH₃OH, we obtain a total column density of (2.0 ± 0.4) × 10¹⁸ cm^-2. Since the A/E abundance ratio is estimated to be 0.77 ± 0.33 (see Sect. 5.1), we simply assume that the A/E abundance ratio is unity for ¹³CH₃OH. We arrive at a total column density of (3.2 ± 0.6) × 10¹⁶ cm^-2 for ¹³CH₃OH. This results in a ¹²C/¹³C ratio of 63 ± 17, which agrees very well with the result (57 ± 14) based on CH₃OH transitions at higher frequencies (Persson et al. 2007) and the result 50 ± 20 from H₂CS transitions (Tercero et al. 2010), while it is slightly higher than the value (43 ± 7) from CN lines (Savage et al. 2002). This confirms that the ¹²C/¹³C ratio in the CR is slightly smaller than the value (77 ± 7) in the local ISM and the solar value (89, Wilson & Rood 1994). Since the solar value represents the local interstellar medium 4.5 Gyr ago, the difference may arise from the Galactic chemical evolution. Isotope selective photodissociation by UV photons can influence the ¹²C/¹³C ratio because of the difference in self-shielding of ¹²C and ¹³C (e.g., van Dishoeck & Black 1988; Savage et al. 2002; Milam et al. 2005). ¹³C is expected to be more easily photodissociated, which will result in a higher ¹²C/¹³C ratio, rather than the lower value that we obtained. This suggests that the effect cannot be significant influencing the ¹²C/¹³C ratio in Orion KL. Previous studies have also shown that chemical fractionation does not play a substantial role in influencing such ratios (e.g., Milam et al. 2005).

5.2.3. ¹⁴N/¹⁵N

Here, we use the average column density of NH₃ and ¹⁵NH₃ to estimate the ¹⁴N/¹⁵N ratio in the HC. For NH₃, we find that the column density obtained from the NH₃ (J = K + 4) transitions is nearly an order of magnitude higher than the values derived from the J = K + 1, J = K + 2, and J = K + 3 transitions, which may be due to different excitation conditions. For NH₃ (J = K + 1, J = K + 2, and J = K + 3) that share similar excitation conditions, we take a sigma-weighted average value of the three groups as the column density of NH₃ in the HC, which is estimated to be (2.4 ± 1.2) × 10¹⁷ cm^-2. For ¹⁵NH₃, we take the average value derived from metastable and non-metastable transitions as the column density of ¹⁵NH₃, which is estimated to be (2.4 ± 0.3) × 10¹⁵ cm^-2. This results in a ¹⁴N/¹⁵N ratio of 100 ± 51, which is roughly consistent with the previous value of 170 $\begin{matrix} + 140 \\ -80 \end{matrix}$ $\hbox{$^{+140}_{-80}$}$ by Hermsen et al. (1985), but less than the 234 ± 47 derived from CN transitions (Adande & Ziurys 2012). This indicates that the ¹⁴N/¹⁵N ratio in the HC is smaller than the 450±22 in the local ISM (Wilson & Rood 1994), the 1000 ± 200 in the prototypical starless cloud core L1544 (Bizzocchi et al. 2013), and the 300 ± 50 in Barnard 1 (Daniel et al. 2013). Fractionation in nitrogen is likely to play an important role in cold temperature (~10 K) regions (Rodgers & Charnley 2008), but the effect can be neglected in Orion KL because of its high temperature. Meanwhile, the ¹⁴N/¹⁵N ratio is found to be 361 ± 141 in the prototypical PDR region Orion Bar (Adande & Ziurys 2012). Thus, isotope selective photodissociation by UV photons should not affect the ¹⁴N/¹⁵N ratio a lot either. Since ¹⁵N may originate in massive stars and is potentially destroyed in lower mass stars (e.g., Henkel et al. 1994; Wilson & Rood 1994; Chin et al. 1999; Wang et al. 2009), such an enhanced ¹⁵N in a massive star-forming regions like Orion KL can be expected. Furthermore, our value is similar to the 111 ± 17 in the Large Magellanic Cloud (LMC) massive star-forming region N113, the ~100 in the LMC star-forming region N159HW, and the value in the central region of NGC 4945 (Chin et al. 1999).

Fig. 8

Radex radiative transfer calculations on the excitation of ¹⁵NH₃. a) The modeled line ratios $\frac{{}^{15}{NH}_{3} (2, 2)}{{}^{15}{NH}_{3} (1, 1)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (2,\,2)}{^{15}{\rm NH}_{3} (1,\,1)}$}$ (red lines) and $\frac{{}^{15}{NH}_{3} (5, 5)}{{}^{15}{NH}_{3} (4, 4)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (5,\,5)}{^{15}{\rm NH}_{3} (4,\,4)}$}$ (blue lines) as a function of n_H₂ and T_kin. The solid red and blue lines represent the calculated $\frac{{}^{15}{NH}_{3} (2, 2)}{{}^{15}{NH}_{3} (1, 1)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (2,\,2)}{^{15}{\rm NH}_{3} (1,\,1)}$}$ and $\frac{{}^{15}{NH}_{3} (5, 5)}{{}^{15}{NH}_{3} (4, 4)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (5,\,5)}{^{15}{\rm NH}_{3} (4,\,4)}$}$ line ratios which are labeled. The dashed lines represent observed line ratios and their upper and lower limits. b) The derived likelihood as a function of n_H₂ and T_kin. Only metastable para–¹⁵NH₃ transitions are included in the modeling. The likelihood scale is indicated by the color bar. The contours represent the likelihoods 0.6, and 0.8. c) Same as Fig. 8a, but the solid red lines represent the line ratios of $\frac{{}^{15}{NH}_{3} (3, 2)}{{}^{15}{NH}_{3} (2, 1)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (3,\,2)}{^{15}{\rm NH}_{3} (2,\,1)}$}$ . d) Same as Fig. 8b, but the non-metastable transitions are also included in the modeling. e) Same as Fig. 8a, but the solid red and blue lines represent the line ratios of $\frac{{}^{15}{NH}_{3} (6, 6)}{{}^{15}{NH}_{3} (3, 3)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (6,\,6)}{^{15}{\rm NH}_{3} (3,\,3)}$}$ and $\frac{{}^{15}{NH}_{3} (4, 3)}{{}^{15}{NH}_{3} (3, 3)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (4,\,3)}{^{15}{\rm NH}_{3} (3,\,3)}$}$ . f) Same as Fig. 8d, but for ortho–¹⁵NH₃ transitions.

5.2.4. D/H

The deuterium fraction can be estimated from the ratio between the column densities of para–NH₃ and para–NH₂D, which have been discussed above and which are given in Table 3. We obtain a D/H ratio of (8.3 ± 4.5) × 10^-3. This is slightly higher than the value 3 × 10^-3 from a previous deuterated ammonia (Walmsley et al. 1987) and a deuterated water study (Neill et al. 2013b), but is consistent with (2 − 8) × 10^-3 from a study of other deuterated molecules (Neill et al. 2013a). This confirms that the D/H ratio for NH₃ is strongly affected by fractionation and is nearly two orders of magnitude higher than the abundance ratio in the interstellar medium and the primordial D/H ratio (~1.5 × 10^-5; e.g. Wilson & Rood 1994; Oliveira et al. 2003; Steigman 2007). Our value is also consistent with the ratio derived from HCN and DCN in the Orion Bar (Leurini et al. 2006; Parise et al. 2009). The enhancement could be attributed to a fossil record of the deuteration of icy dust grain mantles (e.g., Roberts & Millar 2000; Millar 2005) or warm deuterium chemistry driven by CH₂D⁺ (Parise et al. 2009).

Fig. 9

Same as Figs. 8b, d, and f, but with different ¹⁵NH₃ abundances of 8 × 10^-10 and 8 × 10^-8 indicated in the upper left of each panel.

5.3. A RADEX non-LTE model for ¹⁵NH₃

The inversion lines of ammonia are widely used as a Galactic and extragalactic molecular temperature tracer (e.g., Walmsley & Ungerechts 1983; Ho & Townes 1983; Mauersberger et al. 2003). However, the ¹⁴NH₃ transitions tend to become optically thick in high-density regions that contain HCs. Furthermore, Orion KL contains several spatial components that contribute to the ¹⁴NH₃ emission, resulting in difficulties disentangling the different velocity components. In the contrast, the ¹⁵NH₃ transitions can be expected to be optically thin since the ¹⁴N/¹⁵N abundance ratio is around 100 in Orion KL (see Sect. 5.2.3). The transitions of ¹⁵NH₃ only show one velocity component from the HC, so they can be used to constrain the physical properties of the HC. Furthermore, ¹⁵NH₃ has proved to be a good tracer for high-density environments, such as ultra-compact HII regions (e.g., Wyrowski & Walmsley 1996). Here, we present a non-LTE analysis of ¹⁵NH₃ using the RADEX code (van der Tak et al. 2007), where the Einstein A coefficients of ¹⁵NH₃ are from the CDMS database, while its collision rates are taken to be the same as those of ¹⁴NH₃ (Danby et al. 1988). A spherical cloud geometry is assumed. In the models, we do not account for any external radiation fields with the notable exception of the cosmic microwave background. However, infrared pumping is likely to affect the populations of ¹⁵NH₃ in the innermost regions of Orion KL. Therefore, spatial densities derived from purely collisional excitation can be considered as upper limits. To evaluate the models, we compared the observed line ratios with the modeled line ratios to minimize χ², which is defined as $χ^{2} = Σ_{i} \frac{{(R_{obs (i)} - R_{model (i)})}^{2}}{σ_{obs (i)}^{2}},$ $\begin{equation} \label{f:chi} \chi^{2} = \Sigma_{i} \frac{\left(R_{{\rm obs}(i)}-R_{{\rm model}(i)}\right)^{2}}{\sigma_{{\rm obs}(i)}^{2}}, \end{equation}$ (7)where R_obs(i) and R_model(i) represent the observed and modeled line ratios, and σ_obs(i) represents the uncertainty in R_obs(i). We note that the uncertainties of observed line ratios, ( $\frac{{}^{15}{NH}_{3} (2, 2)}{{}^{15}{NH}_{3} (1, 1)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (2,\,2)}{^{15}{\rm NH}_{3} (1,\,1)}$}$ , $\frac{{}^{15}{NH}_{3} (4, 4)}{{}^{15}{NH}_{3} (2, 2)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (4,\,4)}{^{15}{\rm NH}_{3} (2,\,2)}$}$ , $\frac{{}^{15}{NH}_{3} (5, 5)}{{}^{15}{NH}_{3} (4, 4)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (5,\,5)}{^{15}{\rm NH}_{3} (4,\,4)}$}$ , $\frac{{}^{15}{NH}_{3} (3, 2)}{{}^{15}{NH}_{3} (2, 1)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (3,\,2)}{^{15}{\rm NH}_{3} (2,\,1)}$}$ , etc.) should be free of calibration errors and beam dilution effects due to being close in frequency. Following previous χ² studies (e.g., Zhang et al. 2014), we only take solutions with a likelihood L> 0.6 into account, where L is defined as $L = \exp (- χ^{2} / 2) / L_{\max},$ $\begin{equation} \label{f:l} L={\rm exp}(-\chi^{2}/2)/L_{\rm max} , \end{equation}$ (8)where L_max is the maximum likelihood, corresponding to the minimum value of χ².

Assuming that the HC has a typical size of 10′′ and a line width of 7 km s^-1, we obtain a velocity gradient of 350 km s^-1 pc^-1. Adopting a fixed para-¹⁵NH₃ abundance of 8 × 10^-9 (see Table 3), we arrive at a para-¹⁵NH₃ abundance per velocity gradient [X]/(dv/ dr) of 2.3 × 10^-11 pc (km s^-1)^-1. The modeled kinetic temperatures (T_kin) range from 10 to 500 K with a step size of 5 K. The H₂ number density log( $\frac{n (H_{2})}{{cm}^{-3}}$ $\hbox{$\frac{n({\rm H_{2}})}{{\rm cm}^{-3}}$}$ ) varies from 3.0 to 9.0 with a step size of 0.1. We modeled the metastable para-¹⁵NH₃ line ratios, and the results are shown in Figs. 8a and b. We can see that there are two groups of solutions: (1) a warmer group with T_kin = 145 − 190 K and n_H₂ = 10^6.6 − 10^7.1 cm^-3 and (2) a colder group with T_kin = 95 − 150 K and n_H₂ = 10^7.2 − 10^8.0 cm^-3. The first group of solutions is roughly consistent with previous studies (Genzel & Stutzki 1989; Wang et al. 2010; Goddi et al. 2011). The second group of solutions indicates that the ¹⁵NH₃ emission region may have a colder kinetic temperature in an even denser region. This can be expected since the HC in Orion KL is thought to be externally heated (Bally et al. 2011; Zapata et al. 2011; Nissen et al. 2012), which may result in an inner denser region with a lower kinetic temperature. On the other hand, we find that the line ratio $\frac{{}^{15}{NH}_{3} (5, 5)}{{}^{15}{NH}_{3} (4, 4)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (5,\,5)}{^{15}{\rm NH}_{3} (4,\,4)}$}$ is only sensitive to kinetic temperature in a region with a spatial density less than 10⁷ cm^-3. When densities are higher than 10⁷ cm^-3 at the abundance given above, these lines become optically thicker, and thus the ratios are no longer a good kinetic temperature tracer. Including the non-metastable para-¹⁵NH₃ in the modeling (see Figs. 8c and d), we obtain a solution with T_kin = 95 − 145 K and n_H₂ = 10^7.4 − 10^8.0 cm^-3, similar to the second group of solutions based on metastable transitions, which is supportive of a colder and denser region existing in the HC.

We also modeled the line ratios from ortho-¹⁵NH₃. Since the ortho-to-para ratio is found to be around unity (see Sect. 5.1), the ortho-¹⁵NH₃ abundance per velocity gradient is assumed to be the same as for para-¹⁵NH₃. With the line ratios $\frac{{}^{15}{NH}_{3} (6, 6)}{{}^{15}{NH}_{3} (3, 3)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (6,\,6)}{^{15}{\rm NH}_{3} (3,\,3)}$}$ and $\frac{{}^{15}{NH}_{3} (4, 3)}{{}^{15}{NH}_{3} (3, 3)}$ $\hbox{$\frac{^{15}{\rm NH}_{3} (4,\,3)}{^{15}{\rm NH}_{3} (3,\,3)}$}$ , we use the same method to model these ratios from ortho-¹⁵NH₃, and the results are shown in Figs. 8e and 8f. We find that ortho-¹⁵NH₃ emission is likely to come from a region with T_kin = 60 − 105 K and n_H₂ = 10^8.0 − 10^8.6 cm^-3, which is indicative of an even colder and denser region. This further suggests that the HC is externally heated.

The ¹⁵NH₃ abundance calculated with the LTE column density and an assumed H₂ column density may have a large uncertainty, which probably results in a large error in T_kin and n(H₂). Thus, we carried out another RADEX modeling with two different ¹⁵NH₃ abundances of 8 × 10^-10 and 8 × 10^-8. The results are shown in Fig. 9. By comparison with Fig. 8, we find that the derived kinetic temperature depends little on the assumed ¹⁵NH₃ abundances, while the modeled spatial density depends strongly on the assumed ¹⁵NH₃ abundances. Figure 9 also shows that the denser region has a lower kinetic temperature even when changing abundances, therefore, external heating of the HC still holds for the ¹⁵NH₃ abundances ranging from 8 × 10^-10 to 8 × 10^-8.

As mentioned above, an external radiation field is not taken into account in this modeling. In principle, infrared pumping can affect the population of ammonia via vibrationally excited lines (e.g., see Fig. 1 of Mauersberger et al. 1988). When the infrared radiation field becomes very intense, the population can be significantly affected, and this effect may have led to a large number (~20) of NH₃ masers in W51-IRS2 (Henkel et al. 2013). The lack of NH₃ masers and T_kin values below 200 K indicate that vibrational excitation of ammonia does not play a major role for the HC in Orion KL. However, our modeling without excitation by the infrared radiation field only gives upper limits for the spatial density.

6. Summary and conclusions

We have carried out a 1.3 cm spectral line survey toward Orion KL with the Effelsberg-100 m telescope. We detected a total of 261 lines, yielding an average line density of about 23 spectral features per GHz above 3σ (a typical value of 3σ is 15 mJy). Among them, 164 lines are RRLs from hydrogen, helium, and carbon, and 97 lines can be assigned to 13 different molecules including rare isotopologues. A total of 23 molecular transitions from species known to exist in Orion KL were detected for the first time in the interstellar medium. Non-metastable ¹⁵NH₃ transitions were detected in Orion KL for the first time.

Analysis of line widths suggested that the bulk of the discrepancy between observed line widths of hydrogen and helium RRLs stems from their thermal contributions. The intensity ratios of RRL pairs with different Δn from the same atom at neighboring frequencies are consistent with the LTE ratios, suggesting that the LTE deviations are very small in this band. The origin of molecules detected by our survey is discussed according to observed lines and ALMA images. A narrow feature is found in SO₂ (8_1,7 − 7_2,6), but not in other SO₂ transitions, possibly suggesting the presence of a maser line. Column densities and fractional abundances relative to H₂ are estimated for 12 molecules with rotational diagrams.

We found that the four classes of non-metastable NH₃ from J = K + 1 to J = K + 4 have different excitation conditions and metastable ¹⁵NH₃ has a higher excitation temperature than non-metastable ¹⁵NH₃. The elemental and isotopic abundance ratios were calculated: He/H = (8.7 ± 0.7)% derived from the ratios between helium RRLs and hydrogen RRLs; ¹²C/¹³C = 63 ± 17 from ¹²CH₃OH/¹³CH₃OH; ¹⁴N/¹⁵N = 100 ± 51 from ¹⁴NH₃/¹⁵NH₃; D/H = (8.3 ± 4.5) × 10^-3 from NH₂D/NH₃. The dispersion of the He/H ratios derived from Hα/Heα pairs to Hδ/Heδ pairs is very small, which is consistent with theoretical predictions that the departure coefficients b_n factors for hydrogen and helium are nearly identical. A non-LTE model of ¹⁵NH₃ neglecting external radiation fields and a likelihood analysis supports the view that the hot core is externally heated.

Online material

Table 2

Lines detected in the survey of Orion-KL.

Table 4

Spatial origin of molecules detected by our 1.3 cm line survey.

Appendix A: The observed properties of detected lines in the survey

Table A.1

Transitions of recombination lines.

Table A.2

Observed properties of NH₃, ¹⁵NH₃ and NH₂D transitions.

Table A.3

Observed properties of CH₃OH transitions.

Table A.4

Observed properties of SO₂ and OCS transitions.

Table A.5

Observed properties of CH₃OCHO transitions.

Table A.6

Observed properties of H₂O, HDO, CH₃CN, HC₃N, HC₅N, CH₃OCH₃, H₂CO and HNCO transitions.

Appendix B: Zoom-in plots of observed spectra

Fig. B.1

Observed spectrum of Orion KL from 17.9 to 26.2 GHz. The displayed frequency scale is based on the Local Standard of Rest velocity 0 km s^-1.

continued.

continued.

continued.

continued.

continued.

Observed Hα, Heα and Cα transitions indicated by dashed lines with n and Δn (see Sect. 4.1) also given. He63α is blended with H79β while He65α is blended with SO₂ (5_2,4 − 6_1,5). The velocity scale refers to the respective Hα line in each panel.

Fig. B.3

Observed Hβ and Heβ transitions indicated by dashed lines with n and Δn also given. H79β is blended with He63α while H81β is blended with NH₃ (3, 3). The velocity scale refers to the respective Hβ line in each panel.

Fig. B.4

Observed Hγ and Heγ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hγ line in each panel. The spectrum near He100γ is NH₃ (6,2).

Fig. B.5

Observed Hδ and Heδ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hδ line in each panel.

Fig. B.6

Observed Hε and Heε transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hε line in each panel.

Fig. B.7

Observed Hζ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hζ line in each panel.

Fig. B.8

Observed Hη transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hη line in each panel.

Fig. B.9

Observed Hθ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hθ line in each panel.

Fig. B.10

Observed Hι, Hκ and Hλ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hι, Hκ and Hλ lines in each panel.

Fig. B.11

Observed NH₃ transitions with quantum numbers indicated in the upper right of each panel.

Fig. B.11

continued.

Fig. B.12

Observed ¹⁵NH₃ and NH₂D transitions with a one-component Gaussian fit shown (red lines). Quantum numbers are indicated in the upper right of each panel.

Fig. B.12

continued.

Fig. B.13

Observed CH₃OH and ¹³CH₃OH transitions with a one-component Gaussian fit shown (red lines). Species and quantum numbers are given in the upper right of each panel.

Fig. B.13

continued.

Fig. B.14

Observed SO₂ and OCS transitions (black lines) with a two- or three-component Gaussian fit shown together with the individual Gaussian components (red lines). Species and quantum numbers are given in the upper right of each panel. Note that SO₂ (5_2,4–6_1,5) is blended with He65α at 23 413.8 MHz and SO₂ (12_3,9–13_2,12) is blended with He107δ at 20 333.8 MHz.

Fig. B.15

Observed CH₃OCHO transitions with one Gaussian fit shown (red lines). Species and quantum numbers are given in the upper right of each panel. In the CH₃OCHO (2_1,2–1_1,1 E) and CH₃OCHO (2_1,2–1_1,1 A) panels, the blue dashed lines represent the systemic velocities.

Fig. B.16

Observed HC₃N, CH₃CN, HDO, H₂O, HNCO, H₂CO, CH₃OCH₃, HC₅N, and CH₃CH₂CN transitions. Species and quantum numbers are indicated in the upper right of each panel. For HC₃N, CH₃CN, HNCO, and CH₃CH₂CN, the spectra are fitted with the HFS method indicated by red lines. For the HDO and CH₃OCH₃ transitions which are blended, their systemic velocity is indicated by blue dashed lines.

Appendix C: The ALMA maps of detected molecules and comparisons with other studies

We make use of the ALMA line survey archival data to study the distribution of molecules detected by our 1.3 cm line survey. The transitions used are listed in Table C.1. Figure C.1 shows the channel maps of these transitions. Note that the first [0,2] km s^-1 panel of NH₂D (3_2,2s − 3_1,2a) is contaminated by H₂CO (9_1,8 − 9_1,9) at 21 6268.7 MHz.

Recently, Feng et al. (2015) studied Orion KL using combined SMA and IRAM 30 m data, which covers parts of the frequency range of the ALMA-SV data. Here, we make a brief comparison between these data. By comparing the dust continuum emission in both datasets, mm2 in Feng et al. (2015) is resolved into MM4, MM5, and MM6 by ALMA. On the other hand, the dust continuum sources SR, NE, OF1N, and OF1S in Feng et al. (2015) are not detected in the ALMA-SV data. A comparsion of OCS (18–17) demonstrates that Feng et al. (2015) show more extended emission while the ALMA-SV data display more sub-structures. That is because the ALMA-SV data have a higher spatial resolution but lack short spacing and have a smaller primary beam.

We also estimate the flux loss by comparing the ALMA-SV data with those obtained with the IRAM-30 m single dish telescope. We make use of HC₃N (24–23) from Esplugues et al. (2013a), OCS (18–17) from Tercero et al. (2010), and SO₂ (11_5,7 − 12_4,8) from Esplugues et al. (2013b). Based on the telescope information¹⁰, we use a forward efficiency of 94%, a main

beam efficiency of 63%, and a conversion factor from brightness temperature to flux of 7.5 Jy/K to derive the total flux observed by the IRAM-30 m telescope. Meanwhile, the IRAM-30 m telescope has a spatial resolution of ~12′′, covering most emitting regions in Orion KL, so the IRAM-30 m data can be considered to represent the total flux densities of Orion KL. By integrating the whole regions of the ALMA-SV data, we find that the ALMA-SV data can account for 89% of the HC₃N (24–23), 57% of the OCS (18–17), and 44% of the SO₂ (11_5,7 − 4,8) emission.

Table C.1

Transitions from the ALMA line survey..

Fig. C.1

Molecular line channel maps (contours) overlaid on the 230 GHz continuum map (grey). Grey shadings of the continuum image are 10%, 20%, 40%, 60%, 80% of the peak intensity of 1.406 Jy beam^-1. The contour levels of the molecular line images start at 5σ and continue in steps of 5σ, where the σ value for each transition is shown in the first panel in units of Jy beam^-1. The dotted contours are the negative features with the same contour absolute levels as the positive ones in each panel. The symbols are the same as in Fig. 1. The corresponding molecular transitions are indicated in the upper left of the first panel. The velocity range is given in the lower right of each panel in km s^-1. The synthesized beams of the molecular line images are shown in the lower left of each panel. The (0, 0) position in each panel is (α_J2000, δ_J2000) = (05^h35^m14.350^s, −05°22′35.00′′).

continued.

continued.

continued.

continued.

continued.

¹

http://www.cv.nrao.edu~aremijan/PRIMOS/index.html

²

The 100-m telescope at Effelsberg is operated by the Max-Planck-Institut für Radioastronomie (MPIFR) on behalf of the Max-Planck-Gesellschaft (MPG).

³

http://www.iram.fr/IRAMFR/GILDAS

⁴

http://almascience.nrao.edu/alma-data/science-verification

⁵

⁶

http://spec.jpl.nasa.gov

⁷

http://www.astro.uni-koeln.de/cdms/catalog

⁸

http://www.splatalogue.net

⁹

http://www.nist.gov/pml/data/micro/index.cfm

¹⁰

For NH₂D, the inversion motion splits each rotational level (denoted by the quantum numbers J_{k_ak_c}) into two inversion states, labeled as “s” and “a”, corresponding to symmetric and antisymmetric states of the two equivalent minimum-energy structures (Ho & Townes 1983). The ortho levels are those in the “s” state with odd k_a and in the “a” state with even k_a while other states belong to the para species.

¹⁰

http://www.iram.es/IRAMES/mainWiki/Iram30mEfficiencies

Acknowledgments

We would like to thank an anonymous referee for a helpful report that led to improvements in the paper. We greatly thank Christian P. Endres for providing ¹⁵NH₃ spectroscopic information. We wish to thank Zhiyu Zhang for discussions of the χ² analysis, and we appreciate the assistance of the Effelsberg-100m operators during the observations. Y. Gong acknowledges support by the MPG-CAS Joint Doctoral Promotion Program (DPP), and NSFC Grants 11127903, 11233007 and 10973040. S. Thorwirth gratefully acknowledges funding by the Deutsche Forschungsgemeinschaft (DFG) through grant TH 1301/3-2. S. Spezzano wishes to thank the DFG SFB956 and the “Fondazione Angelo della Riccia” for funding. This paper makes use of the following ALMA data: ADS/JAO.ALMA#2011.0.00009.SV. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada) and NSC and ASIAA (Taiwan), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. This research made use of NASA’s Astrophysics Data System.

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All Tables

Table 1

Existing line surveys of Orion KL.

	Fig. 2 Overview of the 1.3 cm line survey toward Orion KL with strong lines marked. The displayed frequency scale is based on the Local Standard of Rest velocity 0 km s^-1.
In the text

	Fig. 4 Comparison of observed and LTE ratios of recombination lines. The observed ratios for hydrogen and helium RRLs are indicated with circles and pentagrams. The red dashed lines represent the LTE ratios, while the blue dashed lines represent the LTE ratios with the departure coefficients b_n corrected. Each ratio is given below the respective abscissa.
In the text

	Fig. 5 Observed peak intensities of CH₃OH (J₂ − J₁E) masers as a function of rotational quantum number J.
In the text

	Fig. 7 He⁺/H⁺ abundance ratios derived from RRLs as a function of rest frequency. The abundance ratios derived from Heα/Hα, Heβ/Hβ, Heγ/Hγ, and Heδ/Hδ are marked with red filled circles, blue filled triangles, green filled squares and black filled diamonds, respectively. The dashed line represents the sigma-weighted mean value of He⁺/H⁺ abundance ratios.
In the text

	Fig. 9 Same as Figs. 8b, d, and f, but with different ¹⁵NH₃ abundances of 8 × 10^-10 and 8 × 10^-8 indicated in the upper left of each panel.
In the text

	Fig. B.1 Observed spectrum of Orion KL from 17.9 to 26.2 GHz. The displayed frequency scale is based on the Local Standard of Rest velocity 0 km s^-1.
In the text

	Fig. B.2 Observed Hα, Heα and Cα transitions indicated by dashed lines with n and Δn (see Sect. 4.1) also given. He63α is blended with H79β while He65α is blended with SO₂ (5_2,4 − 6_1,5). The velocity scale refers to the respective Hα line in each panel.
In the text

	Fig. B.3 Observed Hβ and Heβ transitions indicated by dashed lines with n and Δn also given. H79β is blended with He63α while H81β is blended with NH₃ (3, 3). The velocity scale refers to the respective Hβ line in each panel.
In the text

	Fig. B.4 Observed Hγ and Heγ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hγ line in each panel. The spectrum near He100γ is NH₃ (6,2).
In the text

	Fig. B.5 Observed Hδ and Heδ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hδ line in each panel.
In the text

	Fig. B.6 Observed Hε and Heε transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hε line in each panel.
In the text

	Fig. B.7 Observed Hζ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hζ line in each panel.
In the text

	Fig. B.8 Observed Hη transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hη line in each panel.
In the text

	Fig. B.9 Observed Hθ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hθ line in each panel.
In the text

	Fig. B.10 Observed Hι, Hκ and Hλ transitions indicated by dashed lines with n and Δn also given. The velocity scale refers to the respective Hι, Hκ and Hλ lines in each panel.
In the text

	Fig. B.11 Observed NH₃ transitions with quantum numbers indicated in the upper right of each panel.
In the text

	Fig. B.12 Observed ¹⁵NH₃ and NH₂D transitions with a one-component Gaussian fit shown (red lines). Quantum numbers are indicated in the upper right of each panel.
In the text

	Fig. B.13 Observed CH₃OH and ¹³CH₃OH transitions with a one-component Gaussian fit shown (red lines). Species and quantum numbers are given in the upper right of each panel.
In the text

	Fig. B.14 Observed SO₂ and OCS transitions (black lines) with a two- or three-component Gaussian fit shown together with the individual Gaussian components (red lines). Species and quantum numbers are given in the upper right of each panel. Note that SO₂ (5_2,4–6_1,5) is blended with He65α at 23 413.8 MHz and SO₂ (12_3,9–13_2,12) is blended with He107δ at 20 333.8 MHz.
In the text

	Fig. B.15 Observed CH₃OCHO transitions with one Gaussian fit shown (red lines). Species and quantum numbers are given in the upper right of each panel. In the CH₃OCHO (2_1,2–1_1,1 E) and CH₃OCHO (2_1,2–1_1,1 A) panels, the blue dashed lines represent the systemic velocities.
In the text

A 1.3 cm line survey toward Orion KL⋆,⋆⋆

1. Introduction

2. Observation and data reduction

2.1. Observations with the Effelsberg-100 m telescope

2.2. Archival data

3. Components of Orion KL

4. Results

4.1. Line identifications

4.2. Radio recombination lines

4.3. Molecular lines

4.3.1. NH3, 15NH3, and NH2D

4.3.2. H2O and HDO

4.3.3. Sulfur-bearing molecules, SO2 and OCS

4.3.4. Cyanopolyynes, HC3N and HC5N

4.3.5. CH3OH and 13CH3OH

4.3.6. HNCO

4.3.7. H2CO

4.3.8. CH3CN and CH3CH2CN

4.3.9. CH3OCHO and CH3OCH3

4.3.10. A brief summary

5. Discussion

5.1. Rotational temperatures, column densities and fractional abundances relative to H2

5.2. Elemental and isotopic abundance ratios in Orion KL

5.2.1. He/H

5.2.2. 12C/13C

5.2.3. 14N/15N

5.2.4. D/H

5.3. A RADEX non-LTE model for 15NH3

6. Summary and conclusions

Online material

Appendix A: The observed properties of detected lines in the survey

Appendix B: Zoom-in plots of observed spectra

Appendix C: The ALMA maps of detected molecules and comparisons with other studies

Acknowledgments

References

All Tables

All Figures

A 1.3 cm line survey toward Orion KL^⋆,^⋆⋆

4.3.1. NH₃, ¹⁵NH₃, and NH₂D

4.3.2. H₂O and HDO

4.3.3. Sulfur-bearing molecules, SO₂ and OCS

4.3.4. Cyanopolyynes, HC₃N and HC₅N

4.3.5. CH₃OH and ¹³CH₃OH

4.3.7. H₂CO

4.3.8. CH₃CN and CH₃CH₂CN

4.3.9. CH₃OCHO and CH₃OCH₃

5.1. Rotational temperatures, column densities and fractional abundances relative to H₂

5.2.2. ¹²C/¹³C

5.2.3. ¹⁴N/¹⁵N

5.3. A RADEX non-LTE model for ¹⁵NH₃