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Open Access
Issue
A&A
Volume 698, June 2025
Article Number A47
Number of page(s) 10
Section Stellar atmospheres
DOI https://doi.org/10.1051/0004-6361/202554171
Published online 28 May 2025

© The Authors 2025

Licence Creative CommonsOpen Access article, published by EDP Sciences, under the terms of the Creative Commons Attribution License (https://creativecommons.org/licenses/by/4.0), which permits unrestricted use, distribution, and reproduction in any medium, provided the original work is properly cited.

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1 Introduction

PHOENIX (Hauschildt et al. 1997; Hauschildt & Baron 1999), currently at version 20, is a generalized non-local thermodynamic equilibrium (NLTE) 1D and 3D stellar atmosphere code that has been used for a variety of astrophysical sources including stars (Fuhrmeister et al. 2005, 2010; Passegger et al. 2018), irradiated planets (Barman et al. 2002; Hauschildt et al. 2008; Peacock et al. 2019), novae (Short et al. 2001) supernovae (Mitchell et al. 2002; Baron et al. 2003; Friesen et al. 2017; DerKacy et al. 2020), and AGN (Casebeer et al. 2008). PHOENIX/1D has been used to compute grids of stellar atmosphere models (Barman et al. 2000; Allard & Hauschildt 1995; Hauschildt et al. 1999a,b; Husser et al. 2013) for modeling and analysis of stellar spectra. These model grids have found widespread use in the astronomical community (for example, Smitha et al. 2025; Madore et al. 2025; Savel et al. 2025; Barnes et al. 2024; Passegger et al. 2019; Delgado et al. 2016; Shporer et al. 2014); however, PHOENIX/1D models have substantially improved over the last decade.

Therefore, we present here the NewEra local thermodynamic equilibrium (LTE) grid of spherically symmetric stellar models, computed with the latest version of PHOENIX/1D and the corresponding set of input physical data. An NLTE version of the grid is in preparation and will be discussed in a subsequent paper. As for the previous model grids, the purpose of this grid is to aid in the analyses of stellar spectra, modeling stellar populations, and applications such as transit light curve modeling. The results are available online (see Section 6) in detailed and compact Gaia DR4 formats for different use cases. The NewEra grid covers a similar parameter range as the ACES grid (Husser et al. 2013) and follows the same mass − Teff−log (g) relation.

2 Equation of state: ACES

The Astrophysical Chemical Equilibrium Solver (ACES, Barman et al. 2011) equation of state has been used in PHOENIX beginning with version 16 and the ACES grid (Husser et al. 2013). ACES is a state-of-the-art treatment of the chemical equilibrium in a stellar atmosphere. The method is based upon that of Smith & Missen (1982) incorporating new experimental and theoretical thermodynamic data (Barman et al. 2011) for 839 species (84 elements, 289 ions, 249 molecules, 217 condensates) with all available data updates applied. Chemical equilibrium concentrations for all atomic and molecular species (the species are selected at run time) are determined by the pressure, temperature, and density. This calculation is repeated at each layer in the atmosphere for each model iteration so that the final structure is self-consistent.

3 Line list data

3.1 Atomic lines

For the atomic lines, we use the database from Kurucz (Kurucz 2017, 1992) including all updates up to 2019. This results in a database of about 851 million lines from a total of 389 atoms and ions from H I up to U II (with gaps). The largest contributor by number of lines is Rh V with about 37.5 million lines or 4.4% of the total number of lines, however, this represents a negligible source of opacity for the conditions that we consider in this model grid. For the line opacities, the opacities due to iron-group elements such as Fe, Co, Ni are the most important. In order to dynamically account for the most important opacity sources, we use a line selection procedure that is an evolution of the procedure used in Allard & Hauschildt (1995): for each model and each line in the database its opacity in the line center is compared to the local continuous (bound-free and free-free) opacity at a prescribed number of reference layers in the current (T, Pg) structure. If the ratio of line to continuous opacity is larger than a given parameter (in the NewEra grid this is set to 10−4) for at least one of the reference layers, the line will be included in the calculation, otherwise it will be ignored. In addition, a further test is made so that lines with a ratio of line to continuum opacity larger than a second parameter (in the NewEra grid this is set to 1) for at least one of the reference layers will be considered with individual Voigt profiles, whereas for the weaker lines (line/continuum <1), Gauss profiles are used. With this procedure the total number of atomic lines varies from about 8 × 105 (around Teff = 2300 K) to 2.5 × 106 (around Teff = 5000 K) where about 10% of the selected lines are very strong lines with individual Voigt profiles.

Table 1

Molecular lines and their sources used in the NewEra model grid.

3.2 Molecular lines

The molecular line database is significantly different compared to the Husser et al. (2013) model grid. The vast majority of the molecular lines are from the Exomol database (Tennyson et al. 2016). The full list of all molecular lines and their sources used in this work are given in Tables 1 and 2.

The combined molecular line list contains about 823.8 billion individual lines with transition probabilities >0. In order to reduce the size of the line list and to save processing time, we have created a sub-list by the following procedure: for each of the Exomol species and for each provided wavelength band data file individually, we estimate the total line opacity at infinite temperature per molecule and include all lines that contribute more than 10−3 to 10−4 (depending on species) to the estimated total. This procedure results in a list with about 228.9 billion lines. Furthermore, for the model parameters covered by the grid presented here, in particular for Teff ≥ 2300 K, a number of molecules are not present in the atmospheres. When we omit these lines from the list, we obtain a small list with only about 20.5 billion lines. At a nominal resolution of 25 Å the differences between spectra of a model (Teff = 2300 K, log (g) = 4.5, solar abundances) with the full (823.8 G lines) and the smallest (20.5 G) line lists peak around 4 μm with a maximum flux difference less than 0.4%.

4 Model grid description

The new model grid follows the parameters and setup of Husser et al. (2013), this includes the Teff−log (g) – (mixing length, mass) relations, the base solar abundances used, and the abundance pattern variations. The overall parameter range provided is 2300 K ≤ Teff ≤ 12 000 K (with 100 K steps below 8000 K and 200 K steps above 8000 K), 0.0 ≤ log (g) ≤ 6.0 (with steps of 0.5 dex) and metallicities [M/H] from −4.0 to +0.5 in steps of 0.5 dex. For [M/H] from −2.0 to 0.0 we provide additional α element variations from −0.2 to +1.2 with 0.2 dex steps for a subset of the (Teff, log (g)) parameters of the model grid.

The model grid is not ‘square’ in the sense that for any possible parameter combination there is a model available. For example, low gravity models will become unstable against radiation pressure depending on effective temperature and composition. Such non-static models are not included in this grid. For the solar metallicity ([M/H]=0.0) subset, Figure 1 shows which models are available. Overall, 37438 models are included with the data described below.

Table 2

Molecules used in this work (cont’d).

4.1 Code checks

Computing a model grid on these scales is impractical on local computers and multiple HPC resources were used to compute models and synthetic spectra: Hummel2 (Hamburg), HLRN-4 (Berlin and Göttingen), as well as NERSC (Perlmutter CPU and GPU partitions). Gaps in the model grid (e.g., due to random hardware or I/O problems on the supercomputers) were filled using local resources (M1 Mac Studio, M1 Mac Mini, Linux Xeon PC). Thus, it is very important to verify that these very different systems all produce the same model structure. We have run a number of tests on all systems in order to verify that the results for model iterations agree to better than 10−6 relative accuracy in all quantities saved in the model restart file (e.g., temperatures, gas and electron pressures, densities, opacity averages etc.). These tests included serial and MPI runs (with up to 512 MPI processes), different compilers: gfortran, Intel ifort and ifx, NVIDIA nfortran CPU and GPU (OpenACC and OpenMP/offload), AMD flang and NEC fortran/VE on different CPU/GPU architectures: Apple M1, X86_64, AMD64, NVIDIA A100 GPU, NEC VE. During these tests a number of improvements were implemented (e.g., to make the results of the line selection process completely consistent for different MPI setups) before the model grid was computed.

thumbnail Fig. 1

Available models for solar metallicity in a Hertzsprung-Russell Diagram. The parameters of the models are the effective temperature (on the x-axis), log (g) (indicated next to the series of models for constant log (g)), the metallicity (fixed to [M/H]=0 in this plot) and stellar mass (color coded) following Husser et al. (2013). The luminosity of the models follows from the given parameters. See Section 4 for a discussion of the shape formed by the available models.
thumbnail Fig. 2

Effects of line blanketing on the spectral energy distribution for a NewEra grid model with Teff = 3000 K, log (g) = 5.0, and solar abundances (black line). The structure of the grid model was used to compute the ‘true continuum’, i.e., the SED without any spectral lines included (red line). In addition, the blackbody for the effective temperature is shown in the green curve.

4.2 Effects of line blanketing on the spectra

The effects of the line blanketing on the spectral energy distribution is very large, in particular for models with low effective temperatures. This is shown in Figure 2 for a model with Teff = 3000 K, clearly showing the line blanketing effects for a large wavelength range. The ‘true continuum’ can differ by large factors from the SED with spectral lines included.

4.3 Effects of stellar mass

In the NewEra grid, each model is available for one stellar mass. The effects of varying stellar mass for fixed (Teff, log(g)) are for example discussed in Lester et al. (2017) and Neilson et al. (2022). In Figure 7 we show the effects on the spectral energy distribution of varying the model mass by a factor of 4 for a low gravity NewEra grid model. The radial extension (Rout−Rin)/Rin varies from 20% for 0.5 M⊙ to 16% for 0.93 M⊙ and to 10% for 2.0 M⊙. Depending on the required relative accuracy, models with individualized masses may be needed for low gravities.

4.4 Comparison to previous grids

In Figures 3 and 4, we compare representative spectra in the optical to near-IR for four effective temperatures to up to three previous model generations. All spectra were reduced to a resolution of 25 Å by convolution with a Gaussian for clarity. For Teff = 9800 K and 5000 K the differences between the Husser et al. (2013) and NewEra models are small, mostly due to a different treatment of line dissolution close to ionization edges (hydrogen line dissolution at lower electron pressures calibrated using the results of NLTE modeling of the spectrum of Sirius A, Aufdenberg et al., in prep.). At the two lower effective temperatures the differences are larger. Here the differences between NewEra, Gaia DR1 (Kučinskas et al. 2005) and NextGen (Hauschildt et al. 1999b) are due to vastly different molecular line data, this is visible at all plotted wavelengths. In the near-IR, above about 1.1 μm, at this low resolution the spectra of NewEra and Husser et al. (2013) differ little. This is due to very similar water line data so that the changes at the low resolution are not easily visible. At optical wavelengths the differences are much larger, here the significantly different molecular line data have a large impact on the synthetic spectra. In Figure 5 and 6 we plot the temperature-pressure structures for the different model generations. The differences between the NewEra and ACES structures are small, indicating that the actual flux averaged opacities are quite similar whereas the detailed spectra show significant differences.

4.5 Center-to-limb variation

All models in the NewEra grid have been calculated in spherical symmetry (this is also the case for the previous model grids NextGen, Gaia DR1, and ACES). Although for dwarf models the effects of spherical symmetry on the structure of the atmosphere is small, there are significant effects on the center-to-limb variation (limb darkening, see, for example, Chapman 1966) compared to plane parallel models, as shown in Figure 8. The model calculations are performed on an adaptive computational grid generated so that the outer parts of the model atmosphere are optically thin at all wavelengths. In plane parallel computations the intensities at the limb (μ = 0)1 approach a finite value I(μ = 0) > 0, whereas in spherical symmetry the intensities drop to zero below a wavelength dependent μ threshold as the characteristics (the paths of light) closer to the edge of the star become more and more transparent resulting in large differences compared to plane parallel limb darkening. The μ at which the spherical atmosphere becomes transparent (thus setting the apparent size of the stellar disk) depends on the wavelength due to the variation of the opacities with radius. Thus, the apparent size of the model star is always smaller than the size of the computational grid, that is, the last radial grid point.

This effect is much larger and more complex for giants compared to dwarfs, see Figure 9. Dwarfs with a log (g) = 5 have a much smaller extension of the atmosphere, causing a much better-defined edge of the star compared to the giants with log (g) = 0.0 and 0.5. The larger extension of the giant atmospheres causes both a somewhat slower drop-off in the outer atmosphere and larger electron temperature differences, which in conjunction with the temperature sensitivity of the line opacities causes the apparent size of the star to be more wavelength dependent than is the case for dwarfs. However, even for dwarfs the details of the center-to-limb variation are important for interpreting and modeling, e.g., transit light curves. In applications the PHOENIX/1D limb darkening can match the actual limb darkening in the case of transiting systems (see, for example, Kreidberg et al. 2014, their Extended Data Figure 6).

thumbnail Fig. 3

Grid spectra for model with Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for solar metallicity respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work).

4.6 Data provided

All model grid and synthetic spectrum data are provided as HDF5 (.h5) (The HDF Group 1997–2025) files with one file per model. The models are grouped by metallicity [M/H] (for the sake of brevity and filesystem compatibility, Z stands for [M/H] in filenames) and alpha element variation (if different from zero). Thus, all models with solar abundances are in a subdirectory Z–0.0, models with metallicity 10−4 are in the subdirectory Z–4.0 and so on. Correspondingly, the subdirectory Z–1.0.alpha=0.2 contains models with an alpha element enhancement of 0.2 dex for a base abundance pattern of [M/H]= −1.0 or 1/10 of the solar metallicity.

Each file follows a naming convention that gives the model parameters Teff, log (g),[M/H], and α variation (if ≠ 0) as well as important code setup details such as equation of state and model year. The format has the form lteTTTTTGGGGGZZZZ.alpha=AAAA.PHOENIX-NewEra-ACES-COND-2023.HSR.h5$\begin{align*}\mathtt{lteTTTTTGGGGGZZZZ{.}alpha{=}AAAA{.}PHOENIX}&\\ \mathtt{\hbox{-}NewEra\hbox{-}ACES\hbox{-}COND\hbox{-}2023{.}HSR{.}h5}&\end{align*}$

where lte indicates the LTE part of the model grid, TTTTT is the effective temperature, GGGGG is −log(g), ZZZZ the base metallicity, and AAAA the α element enhancement, e.g., lte04400-5.00-0.0.alpha=-0.2.PHOENIX-NewEra-ACES-COND-2023.HSR.h5$\begin{align*}\mathtt{lte04400\hbox{-}5.00\hbox{-}0.0.alpha{=}\hbox{-}0.2.PHOENIX}&\\ \mathtt{\hbox{-}NewEra\hbox{-}ACES\hbox{-}COND\hbox{-}2023.HSR{.}h5}&\end{align*}$

is a model with Teff = 4400 K, log(g) = +5.00, [M/H]= 0.0, and α = −0.2. The entry PHOENIX-NewEra describes the model generation and ACES-COND the equation of state (in this case ACES) with condensation included, 2023 is the model year and HSR the data product (High Sampling Rate for the wavelengths in the spectrum).

The HSR (high sampling rate) spectra that we provide have an intervals based sampling rate2, λ/Δλ, of at least one million (except from 5.8 μm to 10 μm, i.e., between the infrared M and N bands where the sampling rate is between half a million and one million). The details are given in Table 3. The overall wavelength range (in vacuum) covered is from 900 Å to 30 μm, resulting in a total of about 13.1 million sampled wavelength points. The sampling distances Δλ in the longer wavelength regions are integer multiples of those in the shorter wavelength regions. Therefore, it is possible to create evenly spaced spectra across region borders by dropping data points in the shorter wavelength regions.

It is important to remember that the sampling rate is not the same as spectral resolution. In order to compare to a spectrum of a given resolution, the HSR spectra have to be convolved to the target resolution with an appropriate filter, e.g., with a Gaussian. For meaningful results, the target resolution should be less than about 1/10 of the HSR spectrum sampling rate.

The HSR spectrum is given in the HDF5 group /PHOENIX_SPECTRUM in the h5 files in the following datasets:

nwl:: number of wavelength points

w1:: wavelength array (Å in vacuum)

flux:: flux array log10(Fλ)(erg/s/cm2/cm)

bb:: Planck function array log10(Bλ(Teff))(erg/s/cm2/cm).

For reference a low sampling rate (LSR) spectrum is given in the HDF5 group /PHOENIX_SPECTRUM_LSR in the datasets

wl:: wavelength array (Å in vacuum)

fl:: flux array log10 (Fλ)(erg/s/cm2/cm).

In contrast to the HSR spectra, the LSR spectra cover the whole spectral range from the soft X-ray to the radio wavelength range but at a much lower sampling rate. The details can be found in Table 4.

The PHOENIX/1D input Fortran namelist of the calculation of the HSR spectrum from the pre-calculated model atmosphere is saved in the dataset /PHOENIX_NAMELIST/phoenix_nml. This is included to quickly extract parameters that are not coded in the name of the file, e.g., mass of the star, micro-turbulent speed, or the mixing length (using the provided python routine, see Section 6), however, most of the contents is not intended for data users but is necessary to recreate the models if needed.

The HDF5 dataset /PHOENIX_RESTART/phx_restart is a string representation of the PHOENIX/1D model restart file used to create the spectrum given in the file so that new spectra can be generated that use the exact same model atmosphere structure.

The contents of this dataset should be used to extract structure information, such as radii using the provided python routine, see Section 6. The dataset also includes internal PHOENIX/1D data that are only needed for PHOENIX/1D runs and are included here for easier re-running of models with PHOENIX/1D. The outermost radius is the one that should be used when scaling the flux to that of a star with a different radius or when calculating the flux of the model as observed from a given distance.

The regular (stdout) PHOENIX/1D printout for the HSR spectrum PHOENIX/1D run is given as a string in the HDF5 dataset /PHOENIX_STDOUT/phx_stdout. This information is provided for reference and debugging purposes only.

thumbnail Fig. 4

Grid spectra for models with Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for [M/H] = −3.0 respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work).
thumbnail Fig. 5

Structure of models with (bottom to top) Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for solar metallicity respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work) (not all structures are available at all metallicities).
thumbnail Fig. 6

Structure of models with (bottom to top) Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for [M/H] = −3.0 respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work) (not all structures are available at all metallicities).
thumbnail Fig. 7

Effect of (model) mass on the spectral energy distributions (SED) at a resolution of 25 Å for an example model with Teff = 2500 K, log(g) = 0.0, M = 0.93 M⊙, and solar abundances. Δ F/F is defined as (F(mass =[0.5, 2.0])−F(mass =0.93))/F(mass =0.93).

5 Summary and conclusions

We have calculated an LTE grid of stellar models over a wide range of parameters. The total grid consists of 37438 spectra different synthetic spectra in the temperature range 2300 K ≤ Teff ≤ 12 000 K and 0.0 ≤ log(g) ≤ 6.0 and metallicities [M/H] from −4.0 to +0.5 for all available [α/Fe]. The number of atomic and molecular lines is an order of magnitude larger than in our previous grids. We also supply a Python script to access these data files, which should facilitate their use in studies of, for example the Gaia data set. Future work will involve inclusion of NLTE effects into this grid (Hauschildt et al., in prep.) and extension of this LTE grid to lower temperatures, applicable to planetary atmospheres (T. Barman et al., in prep.).

6 Data availability & software products

In order to facilitate access to the model and spectra, we provide a python code get_NewEra_from_FDR.py to download individual models at DOI 10.25592/uhhfdm. 16722 (Hauschildt et al. 2025)3. The script takes Teff, log(g), [M/H], and the α scale as arguments and downloads the corresponding file. A list of all available models with download links and MD5 checksums is given in list_of_available_NewEra_models.txt.

In addition to the main data product described in the previous section, we provide low resolution spectra in a format compatible to the Gaia spectral library definition. The archive NewEra_for_GAIA_DR4.tar contains the spectra in the spectral range and resolution required for the Gaia mission, whereas the archive PHOENIX-NewEra-LowRes-SPECTRA.tar.gz contains the spectra from 2500 Å to 2.5 μm with a resolution (Gaussian filter standard deviation) and sampling rate of 0.1 Å, and finally the archive PHOENIX-NewEra-JWST-SPECTRA.tar.gz gives the spectra in the wavelength range and resolution useful for JWST: between 0.6 μm and 28.5 μm with a resolution (Gaussian filter standard deviation) and sampling rate of 2 Å. Note that the wavelengths and fluxes in these files are given in nm and W/m2/nm, respectively, in order to be compatible to the Gaia spectral library. Each archive contains the spectra for each abundance pattern in individual files and were generated for the HSR spectra by convolution with a Gaussian filter. example_read_gaia_fmt.py is a simple example python code to read the first spectrum of an individual archive file.

We also provide the Python routine example_read_HSR_H5.py to access the contents of the provided h5 files. The routine reads the supplied HDF5 files and returns the data as strings, scalars and arrays. This includes examples of the stored metadata for the particular grid entry and additional output information that can be accessed. The python routine example_read_structure_from_HSR_H5.py shows how to access the structure data in the /PHOENIX_RESTART/phx_restart dataset, in particular radii, gas and electron pressures, and electron temperature.

Table 3

Wavelength coverage and sampling of the HSR spectra.

Table 4

Wavelength coverage and sampling of the standard low sampling rate (LSR) spectra.

Due to size restrictions, limb darkening data need to be generated and are made available upon request.

thumbnail Fig. 8

Limb darkening for representative giant and dwarf models (as indicated in the panels) for 4 wavelengths. The model with Teff = 3000 K, log(g) = 0.0 has a mass of 1.35 M⊙ and the one with Teff = 3000 K, log(g) = 5.0 has a mass of 0.27 M⊙. For comparison the simplest plane parallel limb darkening law I(μ)/I(μ=1)=35(μ+23)$I(\mu) / I(\mu=1)=\frac{3}{5}\left(\mu+\frac{2}{3}\right)$ is shown as a dotted curve.
thumbnail Fig. 9

Visualizations of the center-to-limb variation of a model with Teff = 3000 K, log(g) = 0.0 (left panels, 1.35 M⊙) and log(g) = 5.0 (right panels, 0.27 M⊙), solar abundances for 4 wavelengths: 2775 Å (top left), 6500 Å (top right), 9000 Å (bottom left), and 1.49 μm (bottom right). The colors indicated by the color bar give log10(I/I(μ = 1)).

Acknowledgements

Data processing in this work was facilitated by GNU parallel (Tange 2022). The model grid calculations presented here were partially performed at the National Energy Research Supercomputer Center (NERSC), which is supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC03-76SF00098. The authors gratefully acknowledge the computing time made available to them on the high-performance computers HLRN-IV at GWDG at the NHR Center NHR@Göttingen and at ZIB at the NHR Center NHR@Berlin. These Centers are jointly supported by the Federal Ministry of Education and Research and the state governments participating in the NHR (www.nhr-verein.de/unsere-partner). Additional computing time was provided by the RRZ computing clusters Hummel and Hummel2. We thank all these institutions for a generous allocation of computer time. PHH and EB are grateful for a NERSC Science Acceleration Program (NESAP) awards from NERSC (2014-2020) which allowed for the efficient use of “Burst Buffer” technology. PHH gratefully acknowledges the support of NVIDIA Corporation with the donation of a Quadro P6000 GPU used in this research. TB acknowledges support by NASA-XRP grant 80NSSC21K0572. EB acknowledges support by NASA grants JWST-GO-2114, JWST-GO-2124, JWST-GO-4436, JWST-GO-4522, JWST-GO-3726, JWST-GO-4217, JWST-GO5057, JWST-GO-5290, JWST-GO-6023, JWST-GO-6582, HST-AR-17555, and 80NSSC20K0538. Support for programs #2114, #2124, #4436, #4522, #3726, #4217, #5057, #5290, #6023, #6582, and #17555 were provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-03127. AS acknowledges support by the DFG grant SCHW 1358/5-1 within the priority program SPP 1992 “Exploring the Diversity of Extrasolar Planets”.

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1

μ = cos (θ) where θ is the angle between the beam of light and the normal to the emitting surface.

2

The sampling distance Δλ used in the calculation of our synthetic spectra is not to be confused with an observational resolution element Δ λobs originating from observing with a finite aperture.

All Tables

Table 1

Molecular lines and their sources used in the NewEra model grid.

In the text
Table 2

Molecules used in this work (cont’d).

In the text
Table 3

Wavelength coverage and sampling of the HSR spectra.

In the text
Table 4

Wavelength coverage and sampling of the standard low sampling rate (LSR) spectra.

In the text

All Figures

thumbnail Fig. 1

Available models for solar metallicity in a Hertzsprung-Russell Diagram. The parameters of the models are the effective temperature (on the x-axis), log (g) (indicated next to the series of models for constant log (g)), the metallicity (fixed to [M/H]=0 in this plot) and stellar mass (color coded) following Husser et al. (2013). The luminosity of the models follows from the given parameters. See Section 4 for a discussion of the shape formed by the available models.
In the text
thumbnail Fig. 2

Effects of line blanketing on the spectral energy distribution for a NewEra grid model with Teff = 3000 K, log (g) = 5.0, and solar abundances (black line). The structure of the grid model was used to compute the ‘true continuum’, i.e., the SED without any spectral lines included (red line). In addition, the blackbody for the effective temperature is shown in the green curve.
In the text
thumbnail Fig. 3

Grid spectra for model with Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for solar metallicity respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work).
In the text
thumbnail Fig. 4

Grid spectra for models with Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for [M/H] = −3.0 respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work).
In the text
thumbnail Fig. 5

Structure of models with (bottom to top) Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for solar metallicity respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work) (not all structures are available at all metallicities).
In the text
thumbnail Fig. 6

Structure of models with (bottom to top) Teff = 2300, 3000, 5000, 9800 K and log(g) = 5.0, 5.0, 4.5, and 4.5 for [M/H] = −3.0 respectively, compared to the NextGen (Hauschildt et al. 1999b), Gaia DR1 (Kučinskas et al. 2005), ACES (Husser et al. 2013), and New Era (this work) (not all structures are available at all metallicities).
In the text
thumbnail Fig. 7

Effect of (model) mass on the spectral energy distributions (SED) at a resolution of 25 Å for an example model with Teff = 2500 K, log(g) = 0.0, M = 0.93 M⊙, and solar abundances. Δ F/F is defined as (F(mass =[0.5, 2.0])−F(mass =0.93))/F(mass =0.93).
In the text
thumbnail Fig. 8

Limb darkening for representative giant and dwarf models (as indicated in the panels) for 4 wavelengths. The model with Teff = 3000 K, log(g) = 0.0 has a mass of 1.35 M⊙ and the one with Teff = 3000 K, log(g) = 5.0 has a mass of 0.27 M⊙. For comparison the simplest plane parallel limb darkening law I(μ)/I(μ=1)=35(μ+23)$I(\mu) / I(\mu=1)=\frac{3}{5}\left(\mu+\frac{2}{3}\right)$ is shown as a dotted curve.
In the text
thumbnail Fig. 9

Visualizations of the center-to-limb variation of a model with Teff = 3000 K, log(g) = 0.0 (left panels, 1.35 M⊙) and log(g) = 5.0 (right panels, 0.27 M⊙), solar abundances for 4 wavelengths: 2775 Å (top left), 6500 Å (top right), 9000 Å (bottom left), and 1.49 μm (bottom right). The colors indicated by the color bar give log10(I/I(μ = 1)).
In the text

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