Flare activity and photospheric analysis of Proxima Centauri · 2017-06-16 · Methods. We identify...

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Astronomy & Astrophysics manuscript no. 054d c©ESO 2017June 16, 2017

Flare activity and photospheric analysis of Proxima Centauri

Y. Pavlenko1, 2, 3 A. Suárez Mascareño4, 5, 7, R. Rebolo4, 5, 6, N. Lodieu4, 5, V. J. S. Béjar4, 5, J.I. González

Hernández4, 5

1 Main Astronomical Observatory of the National Academy of Sciences of Ukraine. e-mail: [email protected] Instituto de Astrofísica de Canarias (IAC), La Laguna, Tenerife, Spain3 Center for Astrophysics Research, University of Hertfordshire, College Lane, Hatfield, Hertfordshire AL10 9AB, UK4 Instituto de Astrofísica de Canarias (IAC), Calle Vía Láctea s/n, E-38200 La Laguna, Tenerife, Spain.5 Departamento de Astrofísica, Universidad de La Laguna (ULL), E-38205 La Laguna, Tenerife, Spain.6 Consejo Superior de Investigaciones Científicas, CSIC, Spain,7 Observatoire Astronomique de l’Université de Genève, E-1290 Versoix, Genève, Switzerland.

June 16, 2017

ABSTRACT

Context. We present the analysis of emission lines in high-resolution optical spectra of the planet-host star ProximaCentauri (Proxima) classified as a M5.5V.Aims. We carry out the detailed analysis of observed spectra to get a better understanding of the physical conditionsof the atmosphere of this star.Methods. We identify the emission lines in a serie series of 147 high-resolution optical spectra of the star at differentlevels of activity and compare them with the synthetic spectra computed over a wide spectral range.Results. Our synthetic spectra computed with the PHOENIX 2900/5.0/0.0 model atmosphere fits pretty well the ob-served optical-to-near-infrared spectral energy distribution. However, modelling strong atomic lines in the blue spectrum(3900–4200Å) requires implementing additional opacity. We show that high temperature layers in Proxima Centauriconsist in at least three emitting parts: a) a stellar chromosphere where numerous emission lines form. We suggest thatsome emission cores of strong absorption lines of metals form there; b) flare regions above the chromosphere, wherehydrogen Balmer lines up to high transition levels (10–2) form; c) a stellar wind component with Vr=−30 kms−1 seenin some Balmer lines as blue shifted emission lines. We believe that the observed He line at 4026Å in emission can beformed in that very hot region.Conclusions. We show, that real structure of the atmosphere of Proxima is rather complicated. The photosphere of thestar is best fit by a normal M5 dwarf spectrum. On the other hand emission lines form in the chromosphere, flareregions and extended hot envelope.

Key words. stars: abundances - stars: atmospheres - stars: individual (Proxima) - stars: late type

1. Introduction

M dwarfs are the most numerous and longest-lived starsin our Milky Way, see Kirkpatrick et al. (2012). Unfortu-nately, the determination of the basic parameters of thesestars is hampered by the complicated physical processestaking place in their atmospheres which limit our ability toreproduce their spectra with synthetic models. Due to thelow temperatures and high pressures in M dwarf photo-spheres, modelling their spectra requires detailed account-ing for molecules when dealing with chemical equilibriumin their atmospheres. M dwarf spectra are governed by ab-sorptions of the numerous band systems of diatomic andpoly-atomic molecules. Spectra of M dwarfs also show emis-sion lines which can be formed only in the outermost hightemperature layers of their atmospheres.

Proxima Centauri (= 2MASS J14294291−6240465;GJ 551, V645 Cen) is the closest red dwarf to the Sunlocated at a distance of 1.3019±0.0018 pc (Lurie et al.2014). Because of its proximity, its angular radius can bemeasured directly via interferometry (Kervella et al. 2003).

Send offprint requests to: Yakiv Pavlenko

Its mass is about an eighth of the Sun’s mass, its luminosityis only 0.15% of that emitted by the Sun, its spectral typeis M5.5 (Bessell 1991), its effective temperature is 3050 K,and its density about 40 times that of our Sun. Since itsdiscovery, Proxima Centauri has been suggested to bethe third component of the α Centauri system. Recently,Kervella et al. (2017) based on new observations claimthat Proxima and α Cen are gravitationally bound with ahigh degree of confidence. It may be the third componentof the Alpha Centauri system with a projected physicalseparation of 15,000±700 au (Wertheimer & Laughlin2006).

Proxima Cen is a known flare star that exhibits randombut significant increases in brightness due to magnetic ac-tivity (Christian et al. 2004). The spectrum of Proxima Cencontains numerous emission lines, see Fuhrmeister et al.(2011). These features most likely originate from plage,spots, or a combination of both. In general, the flare rateof Proxima Cen is lower than that of other flare stars ofsimilar spectral type, but is unusually high given its slowrotation period (Davenport et al. 2016). The star has an es-timated rotation period of ∼ 83 days and a magnetic cycle

Article number, page 1 of 18

http://arxiv.org/abs/1706.04678v1

A&A proofs: manuscript no. 054d

of ∼7 years (Benedict et al. 1998; Suárez Mascareño et al.2015, 2016; Wargelin et al. 2017). The X-ray coronal andchromospheric activity have been studied in detail byFuhrmeister et al. (2011) and Wargelin et al. (2017). Re-cently Thompson et al. (2017) claimed detection of rota-tional modulation of emission lines in the Proxima Censpectrum.

Nowadays M dwarfs represent important targets forsearches of exoplanets, and in particular, rocky plan-ets. Given their small radius and low mass, planets areeasier to detect around M dwarfs because the depthof their transits and the amplitudes of the induced ra-dial velocity variations are larger. First rocky planetswere detected by radial velocity and transits around Mstars (Rivera et al. 2015; Charbonneau et al. 2009). Mostof the rocky planets in the habitable zone have beenfound around these very low-mass stars (Udry et al. 2007;Bonfils et al. 2011; Quintana & Barclay 2014; Torres et al.2015; Wright et al. 2016; Anglada-Escudé et al. 2016).Proxima Cen was recently highlighted as a planet hostmid-M dwarf (Anglada-Escudé et al. 2016). Proxima Cenb orbits its host star with a period of 11.2 days, corre-sponding to a semi-major axis distance of 0.05 AU. Prox-ima Cen b has a mass close to that of the Earth (from 1.10to 1.46 mass of the Earth), with orbit in the temperate zone(Anglada-Escudé et al. 2016).

One may assume that Proxima Cen b is surrounded byan atmosphere with a surface pressure of one bar, implyingthat the planet orbits its host star within the habitable zone(Ribas et al. 2016; Turbet et al. 2016; Garraffo et al. 2016).It is worth noting that classical definition of habitable zonesolely implies the restriction of distance from the centralstar and composition of the planetary atmosphere. How-ever, strong flare activity may move the inner boundary ofthe habitability zone far away from the formally computedpossible radius. For this reason, the detailed characterisa-tion of the flare phenomenon present in the atmosphere ofProxima Cen is of great importance.

In this paper we report on the detection of emissionlines in the optical spectra of Proxima taken with differ-ent instruments ran by the European Southern Observa-tory (ESO). In Section 2, we describe the observations anddata reduction. In Section 3 we analyse the characteristicsof several lines, including Balmer lines, sodium resonancedoublet as well as the Ca II H and K lines. In Section 4, weplace our results into a wider context of activity in low-massstars.

2. Spectroscopic observations

2.1. 3.6-m/HARPS optical spectra of Proxima

We retrieved all the available of Proxima spectra from theHARPS ESO public data archive. The dataset consists in316 spectra collected between June 2004 and May 2016.HARPS (Mayor et al. 2003) is a fibre-fed high resolutionechelle spectrograph installed at the 3.6-m ESO telescopein La Silla Observatory (Chile). The instrument has a re-solving power R∼115 000 over a spectral range from 3780to 6810Å and has been designed to attain very high long-term radial velocity (RV) accuracy. It is contained in a vac-uum vessel to avoid spectral drifts due to temperature andair pressure variations, thus ensuring its stability. HARPSis equipped with its own pipeline providing extracted and

wavelength-calibrated spectra, as well as RV measurementsand other data products such as cross-correlation functionsand their bisector profiles. In order to avoid contaminationof the stellar spectra by the calibration lamp we relied onlyon those spectra taken without simultaneous calibration.The final selection consisted in 147 high resolution spectrataken from the public ESO archive, observed between 2004and 2016 with exposure times ranging from 450 to 1200 s.

For the analysis we use the reduced wavelength-calibrated spectra produced by the HARPS pipeline. Wecorrect every spectrum from the velocity of the star andcreated a high signal-to-noise spectrum by co-adding allthe available spectra.

2.2. Spectra of Proxima in different states of activity

We need very high signal-to-noise spectra in order to per-form a detailed study of the activity processes and spectralfeatures of Proxima. To do so we create two high S/N spec-tra with two different sets of individual spectra. One byco-adding all the available spectra, once set in the barycen-tric frame of reference and corrected from the radial veloc-ity of the star, which gives us a final spectrum containingthe information of the spectral features both in times ofstrong and weak activity of Proxima. We label the result-ing spectrum as ’S’. Then we create a second spectrum forwhich we filter out the spectra obtained during flares, andcreate a mean spectrum representing the times of quiet-ness of the star by once again co-adding the selected spec-tra. We label this spectrum as ’QC’. Flares are identifiedby measuring unusually high levels of Ca II H&K emis-sion and Hα emission. We measure the Mount Wilson Sindex (Noyes et al. 1984) and the Hα index defined byGomes da Silva et al. (2011) following the procedure illus-trated in Suárez Mascareño et al. (2015). Spectra that showa S index of Hα index exceeding the seasonal mean by morethan 3 times the RMS of the whole series are considered inflare state. As a result of the process we obtain two spec-tra with SN > 100 in a wide spectral range in two differentstates (average state and low activity state) which allows usto study the changes in its chromosphere related to changesin its activity level.

2.3. VLT/X-shooter spectra

X-Shooter is a multi wavelength cross–dispersed echellespectrograph (D’Odorico et al. 2006; Vernet et al. 2011)mounted on the Cassegrain focus of the Very Large Tele-scope (VLT) Unit 2. The spectrograph is made of threearms covering simultaneously the ultraviolet (UVB; 3000–5500Å), visible (VIS; 5500–10000Å), and near–infrared(NIR; 10000–24800Å) wavelength ranges thanks to thepresence of two diachronic splitting the light. The spec-trograph is equipped with three detectors: a 4096×2048E2V CCD44-82, a 4096×2048 MIT/LL CCID20, and a2096×2096 Hawaii 2RG for the UVB, VIS, and NIR arms,respectively.

We downloaded public data of Proxima from the Eu-ropean Southern Observatory (ESO) science archive. TheVLT/X-shooter spectra were taken on 15 January 2014 be-tween UT = 6h55 and UT = 7h as part of ESO program092.D-0300. The observation strategy was 2AB cycles of12s, 33s, and 44s in the UVB, VIS, and NIR arms, re-


Pavlenko et al.: Spectrum analysis of Proxima Cen

spectively. The slits of 0.5 arcsec 0.4 arcsec, and 0.4 arc-sec were used, yielding resolving powers of 9900 (3.2 pix-els per full-width-half-maximum), 18200 (92.9 pixels perfull-width-half-maximum), and 10500 (2.2 pixels per full-width-half-maximum) in the UVB, VIS, and NIR arms, re-spectively. The read-out mode was set to 400k and low gainwithout binning.

We reduced the raw dataset with the latest version ofthe X-shooter pipeline (2.8.0)1. The pipeline removes theinstrumental signature to the raw spectra, including biasand flat-field. The spectra are wavelength-calibrated, sky-subtracted and finally flux-calibrated with the associatedspectro-photometric standard star observed as part of theESO calibration plan. The output products include a 2Dspectrum associated with a 1D spectrum. Nonetheless, weextracted the 1D UVB, VIS, and NIR spectra with theapsum task under IRAF2 (Tody 1986, 1993).

3. Results

3.1. Absorption spectra of Proxima

3.1.1. Effective temperature from the X-Shooter’s SED

Rajpurohit et al. (2013) and Passegger et al. (2016) showedthat the effective temperature and gravity of normal fieldM5 dwarfs are Teff = 2900±100 K and log g = 5.0±0.5,respectively. These values are most likely applicable toProxima classified as a M5.5 dwarf with solar metallicity.Such metallicity agrees well with the suggestion that Prox-ima is the third (C) component of the α Cen system, seeReipurth & Mikkola (2012) and references therein.

We fitted our synthetic spectra computed with thePhoenix model atmospheres in the effective temperatureTeff range [2600:3100 K] with incremental steps of 100Kand gravity log g [4.0:5.5] with step 0.5 dex to the ob-served be VLT/X-shooter SED. In our work we adoptedthe ”solar” abundances of Anders & Grevesse (1989), ex-cept for iron abundance log N(Fe) = -4.5 in the scale

∑Ni

= 1.0. These abundances agree with Asplund et al. (2009)within accuracy ∼ 0.1 dex for the most elements. Nev-ertheless, our abundances allow to fit the spectra of theSun and solar like stars in good agreement with other au-thors, using comparative simple 1D model atmospheres,see Ivanyuk et al. (2017). We refer the reader to Pavlenko(2014) and Pavlenko & Schmidt (2015) for a review on in-put data and detailed explanation of the procedure em-ployed to compute the synthetic spectra. Our least squaresfitting procedure is described in Pavlenko et al. (2006b).We choose the best fit of the computed spectra to the ob-served spectrum for the minimum of the S function definedas:

S(fh,∆λ,R) =∑

λ

(Fλ − fh ∗ Fsλ(∆λ,R))

2, (1)

where Fλ and F sλ are the fluxes in the observed and com-

puted fluxes respectively, and fh,∆λ,R are the normaliza-tion flux factor, shift in wavelengths between observed and

1 http://www.eso.org/sci/software/pipelines/2 IRAF is distributed by the National Optical Astronomy Ob-servatories, which are operated by the Association of Universitiesfor Research in Astronomy, Inc., under cooperative agreementwith the National Science Foundation.

computed spectra, instrumental broadening factor R, re-spectively. We created fits for all synthetic spectra fromour grid.

In the left panel of Fig. 1, we show the dependence of Scomputed for the range of adopted parameters in effectivetemperature Teff and gravity log g. We carried out ourminimisation procedure only for the ”good”, i.e. withoutnotable telluric absorption and/or emission features acrossλλ 6650 – 6567 Å, 9300 – 9575 Å, 11082 – 11629 Å, 11889– 11894 Å, 13393 – 15000 Å, 17808 – 19638 Å, see Fig. 1.Shorter wavelengths of λ < 5000 Å were excluded due tosome problems discussed in Section 3.1.5. We find a clearminimum of S for Teff = 2900 K for all considered casesof log g . The dependence on log g is rather weak whenvarying the Teff .

3.1.2. Gravity from absorption lines in NIR VLT/X-shooterspectra of Proxima.

We showed that the fit to observed SED of Proxima sug-gests Teff = 2900 K (Fig. 1). To constrain further the gravityof Proxima, we fit gravity-sensitive absorption lines presentin the optical spectrum. Here we draw main attention tothe profiles of atomic lines. It is worth noting that due tocomplicatedness of spectra of M-dwarfs the comparison ofatomic line profiles in computed and observed spectra is noteasy task. The atomic lines in spectra of M-dwarfs form atthe background of the haze of molecular lines of differentintensity. The molecular features/blends cannot be fittedas good as atomic lines. Nevertheless, comparison of ob-served and computed profiles of atomic lines allows us toconstrain appropriate input parameters. In particular, wefocused here our efforts on the resonance lines of potassiumat 7664.9 and 7698.96Å as well as the subordinate tripletof sodium at 8126Å, which are well-known gravity indica-tors. We show the fits to these lines in Fig. A.1. We canconclude that the optical spectrum of Proxima is best re-produced with solar abundances and log g = 5.0 dex. Ourresults agrees within the uncertainties with the log g = 5.5dex derived by Passegger et al. (2016).

3.1.3. CO bands

In Fig. A.2 we compare the observed spectrum of Proximawith the theoretical spectrum using the ∆v = 2 CO bandscomputed for the 2900/5.0/0 model atmosphere. We cansee that the fit of the synthetic CO bands to the observeddata is a good diagnostic to infer the physical conditionsin M dwarf atmospheres, see Pavlenko (2002). In particularwe find that the synthetic spectra with the proper effectivetemperature match reasonably well the observations andagree with the temperatures derived by empirical methods.We see a rather marginal response to a presence of outer hotatmospheric layers, i.e. of a chromosphere, because CO is avery stable molecule of large dissociation potential (D0 =11.105 eV). In general, the CO bands are seen in absorption.At the resolution of our observations we may conclude alsothat our 12C/13C ratio is consistent with the solar becausethe 13CO bands are weak or even absent in the observedspectrum (Fig.A.2).



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Fig. 1. Left: dependence of S on Teff , log g. Right: the best fit of synthetic spectrum 2900/5.0/0 to the observed VLT/X-shooterspectrum of Proxima. We remove here the spectral regions of the strong telluric absorption in J and K bands, spectral ranges usedfor the fits are indicated by horizontal lines at Fλ = 0.8.

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λ gf E" (eV)log N(Ti)

9638.31 2.44E-01 0.848 -7.05± 0.159647.37 3.68E-02 0.818 -7.05± 0.159675.54 1.57E-01 0.836 -7.35± 0.159688.87 2.45E-02 0.813 -7.35± 0.159705.66 9.79E-02 0.826 -7.20± 0.15

Table 1. Abundances of TiI obtained from the fitsto lines shown in Fig. 2.

3.1.4. Abundances in the Proxima atmosphere

In the previous sections we analysed the saturated lines ofatoms and molecules, which, by definition, show a rathermarginal dependence on the changes of abundances. Wecan assume here that the abundances of the alkali elementsshown in Fig. A.1 do not differ much from the solar case.In the VLT/X-shooter spectra we observe also absorp-tion lines of other elements at the background of the localpseudo-continuum formed by the molecular bands acrossthe optical and NIR spectral ranges. Therefore, to make ourabundance analysis more reliable, we employ a few atomiclines of intermediate strength present in the red part of theoptical spectrum of Proxima where molecular absorption isweaker. We discuss these lines below.

Titanium. Lines of TiI are more numerous than FeI inthe spectrum of Proxima due to the lower potential of ion-isation of titanium. In Fig. 2 we show the fit with the syn-thetic spectrum computed for the 2900/5.0 model atmo-sphere with the solar log N(Ti)=−7.05 (Grevesse & Sauval1998). In Table 1 we list the derived abundances. We find a

Ti abundance of log N(Ti)=-7.20± 0.15 from an average offive TiI lines, suggesting a weak metal deficient atmospherefor Proxima. However, we caution this abundance estimatebecause we see that most of the observed absorption linesare weaker than the lines in the theoretical spectrum com-puted at solar abundance.

Iron. Iron lines are numerous in the spectrum of Prox-ima but not as intense as the TiI lines. Although weak linesare more affected by the uncertainties, the theoretical andobserved spectra agree qualitatively well as shown in theleft panel of Fig. 3. On the right panel of Fig. 3 we dis-play the fit to the observed profile of the intense FeI lineat 8327.06Å (gf = 0.02985), showed also in Fig. A.1. Weobtained good quantitative agreement for log N(Fe)=−4.4dex, similar to the the solar iron abundance within ±0.2dex.

Lithium. The absorption bands of the TiO molecule gov-ern the spectra of M dwarfs around the Li resonance doubletat 6707.8Å. The Li doublet is not seen in the spectrum ofProxima as expected for a fully convective old low-mas star.We compare the observed by HARPS and synthetic spectra



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Fig. 3. Left: fits to the observed FeI and VI lines with the 2900/5.0/0 model atmosphere. Right: fit to the observed FeI line at8327.06Å with log N(Fe)=−4.4

computed with the 2900/5.0/0.0 model atmosphere and theline lists from Plez (1998) and Schwenke (1998). Generallyspeaking, Schwenke’s TiO line list allows better to repro-duce the shape of the blends across spectral range contain-ing Li doublet. However, in our case we can only place anupper limit to the lithium abundance in the atmosphere ofProxima at log N(Li)=−12.04 dex (right panel in Fig. 4).

3.1.5. Atomic and molecular absorption spectra in the3800–4200Å region

Absorption lines of neutral species in the blue spectra of M-dwarfs are expected to be much stronger than in the solarcase due to the lower temperatures and higher pressures inthe regions of the absorption line formation. Moreover, wemay expect resonance lines of neutral metals to appear dueto the changes of their ionisation equilibrium in this lowtemperature regime.

To verify our treatment of the pressure broadening wecompute two spectral regions in the solar spectrum contain-ing rather strong enough lines. We follow the procedure de-scribed in Pavlenko et al. (1995) in our computations. Theprofile of the absorption line is described by a Voigt functionH(a, v) and the damping broadening parameter a changeswith depth in stellar atmosphere. We computed the syn-thetic spectra with the VALD3 line list (Ryabchikova et al.2015; Ryabchikova & Pakhomov 2015) for the 5777/4.44/0solar model atmosphere (Pavlenko 2003) with a micro-turbulent velocity of Vm =1 kms−1 and wavelength steps of0.025Å. In Fig. A.3 we observe a good agreement betweenthe profiles of strong atomic lines in the observed spectrumof the Sun as a star and the computed spectrum. (Fig. A.3).

We computed the synthetic spectrum of Proxima acrossthe 3800–4200Å wavelength range. Molecular absorption isweak or absent at these wavelengths, implying that we cansee deeper layers of the photosphere of Proxima. However,the comparison of the intensities of observed absorptionlines compared with the computations reveals reveals someproblems here:

– A simple analysis of the lines listed in Table 2 shows thatonly lines with low excitation energy (E" < 2 eV) areseen in the observed photospheric spectrum of Proxima(Fig. 5). We see that the lines with higher excitationenergies are weaker (or even absent) in the observedspectrum.

– The strongest atomic lines in the observed spectrumare much stronger in the synthetic spectra computed inthe framework of the classical approach. In other words,

damping pressure effects are more pronounced in thecomputed spectrum where atomic lines have more ex-tended wings.

– Our numerical experiments show that changes of effec-tive temperature by ±200 K or log(g) by −0.5 to −1.0do not improve the fit. We cannot reduce the intensi-ties of saturated lines by reducing the associated abun-dances because we know from the Section 3.1.4 that themetallicity of Proxima is near solar.

We can explain the differences by enhancing the con-tinuum opacity, as shown in Fig. 6 where we compare theobserved spectrum with the newly computed one. To re-duce the strength of resonance lines we could move theline forming region into lower pressure regions of the at-mosphere. In this paper, we use a simple approach sug-gesting χc

ν = χc0ν × κ, where χc0

ν , κ are the conventionalopacity and adjusting parameter, respectively. Enhancingthe continuum opacity across blue spectral range shifts theline forming regions upwards, i.e. to layers of lower pres-sure. As a consequence, the strongest lines in the computedspectrum show weaker wings, in better agreement with ob-servations. We obtain satisfactory fits after implementingadditional continuum opacity in the blue, and can explainthe lack of lines of higher excitation potentials, as that case,the photosphere of the star moves upwards into layers of theatmosphere with lower temperature, where only lines of lowexcitation potential can form. Absorption lines of lower ex-citation energies form above the ”new” photosphere, makingthem less sensitive to these changes (Fig. 5).

3.2. Emission lines

The intensity of the emission features visible in the spec-tra changes during the different activity states of the star.To analyse the difference between the typical activity leveland the more quiet states we created two average spectrarepresenting the typical active state (S) and the quiet state(QC), see section 2.2. The ’QC’ spectrum shows lower in-tensity for most of the measured emission lines, specificallythe intensity of Hα here does not exceed 75% of the maxi-mum emission achieved during strong flares.

3.2.1. Hydrogen Balmer lines

All Balmer lines in the spectrum of Proxima are observedin emission. They show strong variability in time with vari-ations up to a factor of 10 in intensity. As a complementary



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Fig. 4. Left: fit to observed fluxes observed across 6708 Å. Right: fits of synthetic spectra computed line lists of Schwenke (1998)and Plez (1998) for different lithium abundances to the observed 6708Å blend, lithium abundances from the bottom: -11.04, -11.54,-12.04.... Model atmosphere 2900/5.0/0.

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λ Element gf E" (eV)Visibility

4147.67 FeI 7.87E-03 1.485 Yes4149.76 FeI 4.93E-06 0.052 Yes4151.11 Er I 2.73E+00 0.000 Yes4152.17 FeI 5.86E-04 0.958 Yes4153.90 FeI 4.77E-01 3.397 No4154.50 FeI 2.05E-01 2.832 No4154.81 FeI 3.98E-01 3.368 No4156.80 FeI 1.55E-01 2.832 No4159.68 V I 1.86E-02 0.287 Yes

Table 2. Identification of lines in the HARPS spectrum ofProxima shown in the left panel of Fig.5

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Fig. 6. The spectrum of Proxima computed for the model with enhanced continuum opacity across the blue spectral region.

information, we provide movie3 with the time variabilityof the Hα in the series of data obtained with HARPS.We measured the pseudo-equivalent widths (pEW) of theemission lines of the full Balmer serie in the VLT/X-shooterspectrum with the task splot under IRAF, which we reportin Table 3. Interestingly, measurements of pEW of Balmer

3 ftp://ftp.mao.kiev.ua/pub/yp/2017/p/halpha.avi

lines in the averaged HARPS ’S’ spectrum provides valuesof the same order than those reported in the Table.

In Fig. 7 we show the intensity profiles of the Hα linevs Doppler velocity (Vr =∆λ/λ× c), where c is the speedof light. We show the profiles for both the quiet (left; QC)and flare (right; S) states. In our spectra we see the fullBalmer serie, from the Hα line to Hθ, corresponding to the3–2 and 10–2 transitions of the hydrogen atom, respectively.



Line j-i λc pEWÅ Å

Hα 3–2 6562.797±0.02 2.5±0.5Hβ 4–2 4861.323±0.02 5.1±0.2Hγ 5–2 4340.462±0.02 6.5±0.5Hδ 6–2 4101.734±0.02 9.9±0.2Hǫ 7–2 3970.072±0.02 5.4±1.0Hζ 8–2 3889.048±0.02 9.4±0.8Hη 9–2 3835.384±0.02 3.9±0.1Hθ 10–2 3797.898±0.02 4.7±0.4

Table 3. Central wavelengths (λc) and pseudo-equivalentwidths measured on the VLT/X-shooter spectrum and expressedinÅ for the observed Balmer serie

Comparison of their emission profiles provides at least a fewimportant results:

– All Balmer lines show strong variability responding tothe temporal changes of flare activity.

– Self-absorption in the core of the Hα line is observedpractically for all stages of activity of Proxima (seemovie3 in supplementary data). We interpret this phe-nomenon as evidence for the existence of comparativelycool matter outside the flare region. The peak of the lineon the red side is higher than on the the blue, most likelydue to the outward motion of the neutral hydrogen asdiscussed in Fuhrmeister et al. (2011).

– The positions of lines do not change during flare events.We interpret this fact as evidence for a quasi-stationarystate of the region where Balmer emission lines form.

– The profiles of all Balmer lines in emission in the inten-sity vs. doppler velocity parameter space have the sameFHWM despite of the differences in intensity. The simi-larity in the profiles of emission lines displayed in Fig. 7is an evidence that they form in the same region of theatmosphere of Proxima heated by flares.

– Some lines in the Fig. 7 show a well-pronounced peakat Vr = – 30 kms−1 , which we interpret as a indicationof an hot stellar wind moving outwards from the star.This peak is well detected in in Hα and Hβ , and seenas a wide feature in the blue wings of Hγ - Hǫ.

– In the Hζ profile, we again see a well-pronounced emis-sion feature at Vr = –30 kms−1 , however, in this casewe identify this as He I line at 3889.64 Å, see section 4.

– The component shifted blue-ward does not change muchwith the activity phase, suggesting that it probablyforms in the hot ionized plasma flow (stellar wind) farenough from the flare region of the star.

– When the flare activity rises, the increased emission inHα covers the line in the blue wing (see movie3 insupplementary material). This detail is also seen in Fig.14 of Fuhrmeister et al. (2011) at lower level activitytimes.

Photospheric spectrum as well as fluxes show rathermarginal responses on the flare activity. In the following weuse flux ratios measured in VLT/X-shooter spectra in com-bination with measurements of pseudo-equivalent widths ofHα and emission feature in its blue wing in the HARPS’S’ and ’QC’ spectra as shown in Fig. 8.

To determine the Hα flux of Proxima received at Earthwe followed the procedure by Herbig (1985):

fHα= pEWHα

Fλ(6563)

Fλ(5556)Fλ(6563, 0.0)10

−0.4V (2)

where the second term is the average flux ratio at the in-dicated wavelengths for the star. The third term is theVLT/X-shooter flux at 6563 Å, received from a star ofmagnitude V = 0 mag, assumed to be 3.8 10−9 erg cm−2

s−1 Å−1. Despite of the noise in the X-Shooter spectrum at5556Å, we estimated the second term from the X-Shooterspectrum and measured a flux ratio of Fλ(6563)

Fλ(5556)= 3±2.

The pEW values of 2.7± 0.1 and 2.6± 0.1Å were de-termined in ’S’ and ’QC’ spectra, respectively, with the av-erage value of pEWHα

= 2.65 ±0.1 Å. For a magnitude ofV = 11.13 for Proxima (Jao et al. 2014), we derived an av-erage value of fHα

= 1.07 ±0.7 10−12 erg cm−2 s−1.In addition, the second term can be estimated from the

R−I colour, using the relationship found by Hodgkin et al.(1995)

Fλ(6563)

Fλ(5556)= 2.786(R− I)C − 0.932 (3)

where R− I= 2.04 mag for Proxima (Jao et al. 2014), andFλ(6563)Fλ(5556)

= 4.75±0.03, resulting fHα= 1.7±0.1 10−12 erg

cm−2 s−1.As a final result we adopted the mean value of both

measurements, resulting fHα= 1.4±0.4 10−12 erg cm−2 s−1.

From this value of fHα, we obtained a luminosity in the Hα

of LHα= 2.8±0.4 10+26 erg s−1 adopting the distance of

1.30 pc for Proxima from Jao et al. (2014)) and a LHα/Lbol

= 4.5 ±0.4 10−5 adopting the bolometric luminosity of 610+30 erg s−1 from Fuhrmeister et al. (2011).

We determine pEWb of the emission feature seen in theblue emission wing of Hα at Vr = – 30 kms−1 in theaverage HARPS spectrum as shown in Fig. 8, We measurein ’S’ and ’QC’ HARPS spectra pEWb = 0.018 Å, the ratiopEWb /pEWHα

= 0.007. It allows us to estimate the totalenergy emitted by the stellar wind in Hα

LbHα

= 7.0× 10−3 × 2.8× 1026 = 2.0× 10+24. (4)

Assuming complete ionisation of hydrogen in the emit-ting region we determine the number of emitting Hα atoms,i.e lower limit of mass loss:

M = LbHα

/hν∗mp∗Nsec = 3.7∗1019g/yr = 1.8∗10−14M⊙/yr,

(5)

where mp and Nsec are the mass of H I atom and the num-ber of sec in 1 year.

3.2.2. HeI line at 4026.19 and 3888.64 Å

The HeI line at 4026.19Å was identified byFuhrmeister et al. (2011), see their Fig. 13 and Table4. We observe strong variability of this line in the observedspectra. We note that the HeI line has the highest exci-tation energy (E′′ = 20.97 eV (see Table 5) with respectto other emission lines observed across our spectral range.It mostly likely formed in the outermost layers of theatmosphere heated by shock waves. Indeed, the HeI cannotbe formed in the same place where emissions in the coresof absorption lines of neutral metals form. Furthermore,the HeI emission line is broader.

In the ’QC’ dataset, the HeI line shows multiplet struc-ture and looks more intensive in comparison with S state



0

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2

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3

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10-29-28-27-26-25-24-23-2

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He I 3888.64 A

10-29-28-27-26-25-24-23-2

Fig. 7. Profiles of Balmer lines shown in the flux normalised to 1.0 in the center of lines vs. radial velocity obtained in quiet mode(left) and flare mode (right). To simplify the plot profiles are shifted in vertical scale, Hα is shown by solid line. The arrows markthe positions of the Vr = -30 kms−1 component feature and He I 3888.64Å line.

100

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300

400

500

600

700

800

6561 6561.5 6562 6562.5 6563 6563.5 6564 6564.5

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x

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Fig. 8. Scheme of the measurements of the pseudo-equivalentwidths of Hα and Vr = -30 kms−1 feature in HARPS spectraof Proxima.

(Fig 9). We suggest that the strong flares destroy the ex-tended emitted region where the line is formed. Likely,broad component seen in the ’S’ spectrum can be asso-ciated with flare region, where, by definition, dispersion ofvelocities to be larger. In the more quiet modes, we see afew shells moving outwards from the star represented bycomponents equally shifted blue-wards and red-wards sug-gesting a multi-component for the line, see Table 5.

Other He I line of is seen in the blue wing of hydrogen Hζ

line, at 3888.64 Å. Excitation potential of this line is onlya bit lower (19.82 eV) in comparison with 4026Å line, seeTable 4. Both lines do not show any remarkable wavelengthshift, so we may assume they form in in the same extendedquasi stable hot layers.

3.3. Emission cores of resonance lines of atoms and ions

3.3.1. H and K of Ca II

CaII H and K lines are well known as indicators of stellar ac-tivity. These lines are collisional controlled and respond to anincrease in the temperature of the lower chromosphere of quietstars like the Sun or Arcturus (Ayres & Linsky 1975). Hence,the appearance of these lines in the spectrum of Proxima is the

Wavelength Terms j′′–j′ E′′ – E′

(A) (cm−1)

4026.18436 3Po-3D 2–1 169086.87 – 193917.264026.18590 3Po-3D 2–2 169086.87 – 193917.264026.18600 3Po-3D 2–3 169086.87 – 193917.254026.19675 3Po-3D 1–1 169086.95 – 193917.264026.19829 3Po-3D 1–2 169086.95 – 193917.264026.35695 3Po-3D 0–1 169087.93 – 193917.26

3888.60467 3S-3Po 1–0 159856.08 - 185564.963888.64560 3S-3Po 1–1 159856.08 - 185564.693888.64893 3S-3Po 1–2 159856.08 - 185564.67

Table 4. Spectroscopic parameters of the emission HeI lines

NN ∆λ Vr

(Å) (kms−1 )1 −0.560 −41.732 −0.420 −31.293 −0.240 −17.884 −0.080 −4.595 0.090 6.716 0.220 16.397 0.380 28.32

Table 5. Observed components of the emission HeI 4026Å line

evidence of chromosphere. Moreover, these lines are extremelystrong:

– Emission lines of H and K CaII fill the broad absorption linesseen in less active M dwarfs of the same spectral type, suchas GJ699 (Barnard star) as shown in Fig. 10.

– These emission lines do not show any wavelength shift withrespect to the photospheric lines.

– The emission profiles of the H and K lines are more intensein the quiet states, in contrary to what we see for Balmerlines of hydrogen. Likely, big flares might affect the regionsof their formation.



100

200

300

400

500

600

700

800

4025.5 4026 4026.5 4027 4027.5

Nor

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ised

Flu

x

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Proxima, QC S

Fig. 9. Profiles of HeI 4026Å line observed in ’S’ and ’QC’ statesof flare activity

– The self-absorptions seen in the cores of the CaII H andK lines indicates that the temperature drops at the upperboundary of their formation region.

– The intensity of their components changes in time and withactivity, as demonstrated in the movie4 in the supplementarymaterial.

3.3.2. Na H and K resonance lines

Strong emission H and K resonance lines of sodium are notablefeatures in the observed spectra of Proxima. These lines arecontrolled by photoelectric processes (Thomas 1959). For theselines, the ratio of the photo-ionisation sink to the collisionalsink is much larger than 1, see Athay (1972). Therefore, theselines, like the hydrogen lines, show a rather marginal responseto the temperature gradients present in chromospheres. We referthe reader to the response of lithium line on chromospheric-likestructures in the atmosphere of M dwarfs described in Pavlenko(1998b) . Only extreme cases of chromospheric activity can pro-vide emission cores in the photo-electrically controlled lines,which is most likely the case of Proxima.

In Fig. 11, we compared the profiles of K and H lines of theNaI resonance doublet, which indicate a few non-trivial results:

– The strong emission cores seen in the NaI resonance doubletshow significant changes of total emitted energy, see movie5.

– However dispersion of velocities in line forming region changevery little with the activity level. In the more quiet mode(QC), emission profiles of both components are a bit nar-rower, while flares increase the width of the emitted lines.Likely, resonance doublet of NaI forms in the chromosphereof the star.

– The profiles of strong photospheric absorption lines prac-tically do not show any response on the level of activity,suggesting that photospheric layers are not bound or weaklybound with regions governed by stellar activity processes.

3.3.3. Other emission lines

In Table A.1 we provide the list of other emission lines withemission cores seen in the spectrum of Proxima. We find

4 ftp://ftp.mao.kiev.ua/pub/yp/2017/p/cahk.avi5 ftp://ftp.mao.kiev.ua/pub/yp/2017/p/na5890.avi

emission lines of many different elements, from HeI (z=2) toDy (z=66). Our list is more complete than the one given inFuhrmeister et al. (2011) because of the higher spectral resolu-tion of our spectra. Moreover, our list contains both pure chro-mospheric lines and lines of very high excitation like HeI, whichcan be formed in areas governed by shock waves (i.e. formed out-side the chromosphere). We can distinguish them by the widthsof their observed profiles because chromospheric lines formedin the cores of absorption lines are narrower than the lines ofhydrogen and helium formed in the layers with larger velocitydispersion.

We assign some specific labels to some of the lines, as follows:– ’ec’ – absorption lines with emission cores,

– ’ecs’ – emission cores with self-absorption as shown in theleft panel of Fig. 12,

– P Cyg and iP Cyg - absorption lines with emission compo-nents in the red and blue wings, respectively. An inverse P Cygline is shown in the right panel of Fig. 12.

We generally find that most of the strong resonance linesshow emission cores and are often shifted with respect to thecentral wavelengths. Observed emission cores do not form in thespherically symmetrical and stable atmosphere. Moreover, thestrong temporal changes of the emission lines provide evidencethat the layers of the chromosphere where they form are stronglyaffected by the flaring processes. Some other cases show P Cygor inverse P Cyg profiles. These phenomena likely reflect thecomplicate dynamical processes occurring in the atmosphere ofProxima.

4. Discussion on the atmosphere of Proxima

4.1. Photosphere

Analysis of the VLT/X-shooter spectra obtained with interme-diate resolution showed that the atmosphere of Proxima in theoptical and near IR spectral range is similar to the rest of nor-mal M-dwarfs of the same spectral classes M5-M6. The opticalspectral range is governed by Ti and VO bands, molecular bandsof water dominates in the near IR. Our theoretical spectra com-puted for a model atmosphere of Teff = 2900 K fits well enoughthe observed spectral energy distribution in accordance with re-sults of Rajpurohit et al. (2013) and Passegger et al. (2016). Insummary, we obtain rather good fits of our synthetic spectracomputed for the canonical PHOENIX model atmosphere of so-lar metallicity across all spectral ranges observed by VLT/X-shooter , except for slightly lower fluxes in the spectral regionof the K-band, which seems to be a common problem of similarmid M-dwarf investigations, see Pavlenko et al. (2006a).

Our computed spectral energy distributions depend on log grather marginally, still we note a weak trend toward lower log gin accordance with log g = 4.5 found by Mann et al. (2015). Onthe other hand, strong alkali lines provide clear response on logg changes. We obtained better fits to profiles of some observedabsorption lines of Na, Rb, subordinate triplet of Na at 8200Åwith log g = 5.0, see also Rajpurohit et al. (2013), instead oflog g = 5.5, as obtained by Passegger et al. (2016). Our resultis in agreement with the log g = 5.23± 0.14 derived from theinterferometric measurement of the Proxima radius and mass-radius relations (Demory et al. 2009).

The opacity of TiO bands around 4200Å starts to decreasebluewards (Pavlenko 2014). At shorter wavelengths, we exploredeeper regions of the atmosphere of Proxima, where the observedspectrum is dominated by absorption lines of neutral atoms. Theresonance lines of neutral species become very strong. Theoret-ical spectra provide too strong wings due to the effects of pres-sure damping which broadens these lines. In the case of Proximawe obtain the best solution for κ = 80. This value looks veryhigh, but we work in the regime of low temperatures where theconventional opacity is extremely low in this case due to thelow density of the H− ion, which is the main source of opacity



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K(Ca II) H(Ca II)

Hε(H I)

G551GJ699

Fig. 10. Left: emission lines in the spectral region of the Ca II H and K lines in of the star Proxima and and GJ 699, see section3.3.2. Right: the same plot shown over a wider wavelength range. The GJ699 spectrum shown here was downloaded from theHARPS archive.

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Fig. 11. Left: resonance doublet of NaI in different levels of activity. Right: resonance doublet of AlI in comparison with GJ 699,both shown in quiet state.

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Fe I Co IFe I

Dy I

K I

Y I

Sc I

G551 (S)QC

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Fe IFe I

Ti I

Proxima (S)QC

Fig. 12. Left: absorption lines of the neutral metals with emission cores. Left arrow at 4045.81Å marks the self-absorption in thecore of the strong Fe I absorption line. In the ’S’ spectrum we note the strong emission detail in the core of the nearby Dy I λ4045.97Å line is marked by the right arrow. In the ’QC’ spectrum the detail is ”eaten” by self-absorption. Self-absorption in theweak emission core of K I line at λ 4047.21 is seen as well. Right: example of iP Cyg profiles.

in the atmospheres of M stars. In other words, we should findlarger electron densities to increase opacity in the low photo-spheric layers. We suggest that this effect may be created bythe over-ionisation of alkali metals which are the main donors offree electrons. Indeed, photons with λ < 4200Å can ionise theneutral atoms in late-type atmosphere due to the low electrondensities where processes of recombination are not so effective.

The additional opacity affects only the emitted spectrum in theUV and blue wavelengths, occurring deep in the atmosphere.The temperature structure of the upper layers of the model at-mosphere are mainly determined by molecular opacities, whichexplain why we can reliably fit the spectral energy distributionof Proxima at longer wavelengths.



On the other hand, for some spectral diagnostics, it couldalso be that the analysis is too simplified and 1-D LTE synthesesmay indeed underestimate some of the opacity (lines) in theblue. More complex analysis involving 3-D atmospheres fromhydro simulations sometimes find that the missing opacity iseven larger in 3D NLTE models than in 1-D LTE. There arehints in the paper by Fuhrmeister et al. (2011) that much moresophisticated 3-D atmospheres should be applied in the case ofProxima.

4.2. Chromosphere

The chromosphere represents the extended part of the atmo-sphere of Proxima, where narrow emission lines of differentintensities vary with time. We observe temporal variations ofthese lines because of interactions between the flares and thechromosphere. We know that solar flares originate in highchromosphere-corona regions. The strongest flares move down-wards to get deeper into the photosphere. However, our moviesshow mainly temporal changes of intensity of chromosphericlines, while the photospheric spectrum shows marginal responseto flares. In the most cases, even the strongest flares seem to oc-cur in local regions far above the surface/photosphere of Prox-ima. The TiO line forming region is separated from the hot flareregions by the mantle of the cool plasma. Nevertheless, the chro-mosphere of Proxima is very powerful because it shows manyabsorption lines of neutral metals with narrow emission cores.These emission cores most likely form in the chromosphere,above the temperature minimum. The presence of emission coreswithin the lines of neutral alkalies and neutral atoms which arefully controlled by radiation can only occur in the presence ofsteep temperature gradients in the chromosphere, see e.g. NLTEsimulations of emission cores of Li in Pavlenko (1998a).

4.3. Flares

Hydrogen Balmer lines in emission is a common phenomenonin the Sun. Balmer lines are controlled by photo-electrical pro-cesses and are not sensitive to the temperature structure of theatmosphere. Their emission is mainly the result of ionization andcollisional processes created by shock waves following flare event.The Hydrogen emission lines seen in the Proxima spectrum formmainly in flare regions. They are broader than chromosphericlines of CaII, NaI and other neutral metals. This effect mostlikely results from a larger scatter in their velocity distribution,as seen in Fig. 10, where we compare the CaII and Hǫ lines: theCaII is much stronger, but Hǫ is broader.

4.4. Stellar wind

We observe a quasi-stationary component in the blue wing ofthe Hα and Hβ emission in the HARPS spectra. We interpretits presence as a result of the flow of highly ionized plasma witha velocity of Vr =−30 kms−1 . We assume that the observedcomponent most likely relates to the the hot stellar wind outflowgenerated by the high level of flare activity in Proxima. Fromour estimate of the total energy emitted in the Hα line, we infera lower limit of the mass loss of M = 1.8×10−14M⊙/yr becausewe do not consider emitted energies by other emission lines.

Wood et al. (2000, 2001) estimated a four or six times lowermass loss of M ≤ 0.2M⊙ = 0.5 10−14M⊙/yr or 0.3 10−14M⊙/yr

for the cases of average solar wind mass loss M⊙ = 1.3741012 g/sec = 2.3 10−14M⊙/yr (Hundhausen 1997) and M⊙ ∼

1.3 10−14M⊙/yr (Goldstein et al. 1996), respectively. Neverthe-less, measurements by Wood et al. (2000, 2001) are based onthe analysis of the Lα absorption which relates to absorptionby neutral hydrogen atoms. Both estimates relate to differentparts of the stellar wind. It is worth noting that the level of ac-

tivity and, respectively, mass loss for the late type stars shouldchange in time. Indeed, we observe different levels of activity inthe M-dwarfs population of our Galaxy. The spectrum of GJ699used in our paper for the comparison with Proxima providesclear evidences of much lower level of activity. Likely, nowadaysProxima passes its evolutionary epoch of high activity.

In our spectra we see manifestations of cool and hot com-ponents of the stellar wind from Proxima. Cool neutral hydro-gen located above the flare region provides the asymmetricalself-absorption in Hα core, as discussed in Fuhrmeister et al.(2011). We defer a more detailed analysis of this phenomenonto a future paper.

5. Conclusions

In the framework our work at least a few results were obtained:• From the fits spectral energy distributions observed in theoptical and near infrared spectral ranges we obtained effectivetemperature of Proxima Teff = 2900± 100 K.• Fit to profiles of strong atomic lines atomic lines observed inthe optical spectrum of Proxima provides good restriction forthe gravity in atmosphere log g = 5.0± 0.25.• From the analysis of strong resonance and subordinate linesof Na, K, Rb as well as lines of intermediate strengths of Ti Iand Fe I lines formed at the background of TiO and VO bandswe obtain solar abundances of these elements in the atmosphereof Proxima.• From the fits to the observed spectrum across Li resonancedoublet we determined the upper limit of Li abundance logN(Li) = −12.04, which is consistent with the expected depletionof an old fully convective low-mass star.• Photospheric lines observed in the optical and infrared spectraof Proxima can be fitted by standard synthetic models. Howeverobserved strong lines of low excitation potentials in the bluespectral region show narrower profiles than expected indicatinga formation in lower pressure layers in the atmosphere. Wewere able to reproduce their profiles incorporating additionalopacity, which shifts their formation layers upwards, in thelower pressure regions.• In spite of a comparatively high level of activity we foundthat the photospheric spectrum show rather marginal responseon the flare activity of Proxima, except for very strong flares.• The emission lines of hydrogen are good indicators of stellaractivity. At the times of strong flares they become moreintense. On the contrary, strong emission lines formed inthe chromosphere, i.e., H&K Ca II, H&K Na I, reduce theirintensity at the occurrence of strong flares. Likely strong flareschange the structure of chromosphere. These chromosphericoriginated lines are narrower in comparison with hydrogen linesformed in the flare region.• The He I line at 4026.19Å is observed in emission and reducesits intensity in the presence of strong flares,too. Due to thelarger excitation potential of the He I line, it should form inhotter layers above the hydrogen lines formation region. Inthe absence of strong flares the He I emission line at 4026.19Åshows multicomponent profile. Likely, it reflects complicatestructure of the line forming region. In the ’S’ spectrum theintensity of the He I lines is lower than in QC, it means thatthe He I lines formation layers are affected by flares.• In the blue wing of emission hydrogen Hα and Hβ lines wefound an emission component shifted to Vr = 30 kms−1 withrespect to their cores. We interpret the emission componentsas evidence of the hot stellar wind from Proxima. Using thesimple model of complete ionisation of hydrogen atoms in thestellar wind we estimate a minimum mass loss of M = 1.810−14M ⊙ /yr.

Acknowledgements. Based on observations collected at the EuropeanOrganisation for Astronomical Research in the Southern Hemisphere



under ESO programme(s) 087.D-0300(A). This is research has madeuse of the services of the ESO Science Archive Facility. YP thanksfinancial support from the Fundación Jesús Serra for a 2 monthstay (Sept–Oct 2016) as a visiting professor at the Instituto de As-trofísica de Canarias (IAC) in Tenerife. NL and VJSB are supportedby the AYA2015-69350-C3-2-P program from Spanish Ministry ofEconomy and Competitiveness (MINECO). J.I.G.H. acknowledges fi-nancial support from the Spanish MINECO under the 2013 Ramóny Cajal program MINECO RYC-2013-14875, and A.S.M., J.I.G.H.,and R.R.L. also acknowledge financial support from the Spanish min-istry project MINECO AYA2014-56359-P. The authors kindly thankM.A. Bautista, S.N. Nahar, M.J. Seaton, D.A. Verner who supplieddata compiled in the NIST database. This research has made use ofthe Simbad and Vizier databases, operated at the Centre de DonnéesAstronomiques de Strasbourg (CDS), and of NASA’s AstrophysicsData System Bibliographic Services (ADS). We thank the anonymousreferee for his/her thorough review and highly appreciate the com-ments and suggestions, which significantly contributed to improvingthe quality of the publication.

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Appendix A:



0

0.1

0.2

0.3

0.4

0.5

0.6

7620 7640 7660 7680 7700 7720 7740

Nor

mal

ised

Flu

x

Wavelength(A)

Proxima2900/4.5/02900/5.0/02900/5.5/0

0

0.1

0.2

0.3

0.4

0.5

0.6

0.7

0.8

0.9

8100 8150 8200 8250 8300

Nor

mal

ised

Flu

x

Wavelength(A)

Proxima2900/4.5/0.2900/5.0/0.2900/5.5/0.

0

0.1

0.2

0.3

0.4

0.5

7790 7795 7800 7805 7810

Nor

mal

ised

Flu

x

Wavelength(A)

Rb IRb I

Proxima2900/4.5/0.2900/5.0/0.2900/5.5/0.

0

0.1

0.2

0.3

0.4

0.5

0.6

7940 7945 7950 7955

Nor

mal

ised

Flu

x

Wavelength(A)

Rb IRb I

Proxima2900/4.5/0.2900/5.0/0.2900/5.5/0.

0.1

0.2

0.3

0.4

0.5

0.6

0.7

0.8

0.9

8320 8340 8360 8380 8400 8420

Nor

mal

ised

Flu

x

Wavelength(A)

Fe I

Ti I Ti ITi I

Fe I

TiITi I

2900/5.0/0.Proxima

0.25

0.3

0.35

0.4

0.45

22000 22050 22100 22150

Nor

mal

ised

Flu

x

Wavelength(A)

Proxima2900/4.5/02900/5.0/02900/5.5/0

Fig. A.1. Top panel: fits to the observed KI resonance doublet (left), subordinate triplet of NaI (right). it Middle panel: fits toRbI resonance doublet lines. Bottom panel: fits to the observed profiles of TiI, FeI, and NaI lines.



0.2

0.25

0.3

0.35

0.4

0.45

22800 23000 23200 23400 23600 23800

Nor

mal

ised

Flu

x

Wavelength(A)

Proxima2900/4.5/02900/5.0/02900/5.5/0

0.2

0.25

0.3

0.35

0.4

23200 23300 23400 23500 23600 23700

Nor

mal

ised

Flu

x

Wavelength(A)

Proxima2900/5.0/0

13CO

Fig. A.2. Top: fit of our synthetic spectra governed mainly by 12CO bands to the observed VLT/X-shooter spectrum. Bottom:13CO bands in the theoretical spectrum computed for the parameters of Proxima.



0

0.2

0.4

0.6

0.8

1

1.2

3920 3930 3940 3950 3960 3970 3980 3990

Nor

mal

ized

Flu

x

Wavelength (A)

Fe IFe I

Ca II Ca II

Na I

Na I

the Sun5777/4.44/0

0

0.2

0.4

0.6

0.8

1

1.2

3565 3570 3575 3580 3585 3590

Nor

mal

ised

Flu

x

Wavelength(A)

Fe I

Fe I + Ni I Fe I

Fe I + Cr I Fe I + Cr I Fe I Fe I

the Sun5777/4.44/0.

Fig. A.3. Comparison of the strong lines in the spectrum of the Sun as a star Kurucz et al. (1984) and the theoretical spectrumcomputed for 1D model atmosphere of Teff /log g /[Fe/H] = 5777/4.44/0.0.



Table A.1. Emission lines in Proxima spectrum, see section3.3.3

λ Element gf E" (eV) Remarks

3819.57 Cr I 1.16E+00 2.708 S,QC ecs3824.44 FeI 4.35E-02 0.0 S,QC ec3829.36 Mg I 5.93E-01 2.709 S,QC ec3832.30 Mg I 1.33E+00 2.712 S,QC ec3837.60 S,QC3838.29 Mg I 2.50E+00 2.717 S,QC ec3840.75 V I 7.28E-01 0.040 S,QC ecs3847.33 V I 8.51E-02 0.017 S, QC ecs, P Cyg?3853.88 S,QC3856.37 FeI 5.18E-02 0.052 S,QC ec3859.91 FeI 1.95E-01 0.000 S,QC ec3864.10 Mo I 9.77E-01 0.000 S ec, iP Cyg?3870.91 S>QC3872.50 FeI 1.18E-01 0.990 S,QC ec3878.57 FeI 4.18E-02 0.087 S,QC ec3886.28 FeI 8.39E-02 0.052 S,QC3887.05 FeI 7.18E-02 0.915 S,QC3888.66 HeI S,QC3894.03 Cr I 2.24E-02 0.961 S,QC ec3894.08 Co I 1.26E+00 1.049 S,QC ec3894.98 Co I 3.98E-02 0.629 S,QC ec3895.66 FeI 2.14E-02 0.110 S,QC ec3897.88 FeI 1.84E-01 2.692 S,QC iP Cyg3899.71 FeI 2.94E-02 0.087 S,QC ecs3903.16 Cr I 5.89E-03 0.968 S,QC ec3903.90 FeI 1.57E-01 2.990 QC3905.52 Si I 9.10E-02 1.909 S,QC3906.48 FeI 5.71E-03 0.110 S,QC ecs3907.49 sr I 2.28E+00 0.000 S,QC ec3907.93 FeI 7.64E-02 2.759 S,QC ec3908.76 Cr I 8.91E-02 1.004 S ec shifted to red3909.86 V I, Co I 7.94E-02 0.069 S,QC complicate blend3920.26 FeI 1.80E-02 0.121 S,QC ecs shifted to red3922.91 FeI 2.23E-02 0.052 S,QC3926.82 S,QC esc3927.92 FeI 3.01E-02 0.110 S,QC ec3930.30 FeI 3.23E-02 0.087 S,QC ec3939.26 S iP Pyg3940.88 FeI 2.51E-03 0.958 S,QC ec shifted to blue3941.49 Cr I 4.07E-02 1.030 S,QC ec shifted to blue3941.73 Co I 9.33E-03 0.432 QC ecs3944.01 Al I 2.38E-01 0.000 QC,S ecs3948.67 TiI 3.98E-01 0.000 S,QC ec3960.05 S3961.52 Al I 4.75E-01 0.014 QC,S ecs3983.13 QC,S4001.66 FeI 1.26E-02 2.176 QC,S ec blue shifted4021.87 FeI 1.87E-01 2.759 QC,S ec4026.19 QC,S He II4030.75 Mn I 3.21E-01 0.0 QC,S ecs4032.98 Ga I 2.36E-01 0.0 S4033.06 Mn I 2.27E-01 0.0 QC,S ecs4034.48 Mn I 1.44E-01 0.0 QC,S ecC4042.93 QC,S4044.14 19.00 1.20E-02 0.000 QC,S P Cyg4045.97 Dy I 7.08E+00 0.0 S ec4047.21 K I 6.03E-03 0.0 QC,S ecs4052.96 QC,S4062.44 FeI 1.37E-01 2.845 QC,S ecs4063.59 FeI 1.15E+00 1.557 QC ecs4064.21 TiI 1.20E-01 1.053 QC,S ec4071.74 FeI 9.51E-01 1.608 QC,S P Cyg4077.36 Y I 1.89E+00 0.0 QC,S iP Cyg4077.71 Sr II 1.47E+00 0.0 QC,S P Cyg4103.94 QC,S



4111.14 na? P Cyg, ecs4115.18 V I 1.18E+00 0.287 QC,S ec4116.47 V I 4.90E-01 0.275 QC,S ec4116.56 V I 1.47E-01 0.262 QC,S ec4131.99 V I 8.51E-01 0.287 QC,S −30 kms−1 emission. iP Cyg.4132.06 FeI 2.11E-01 1.608 QC,S the same4134.48 V I 5.94E-01 0.301 QC,S ecs

4158.62 ? QC,S O II?4158.67 ? QC,S O II?4159.68 V I 1.86E-02 0.287 QC,S ec redshifted.4164.658 Nb I 7.413E-01 0.049 QC,S in -30 kms−1 emission line4167.270 Gd I 1.542E-01 0.124 QC,S P Cyg?4169.877 Ce II 4.467E-01 0.536 QC,S P Cyg?4173.44 QC,S iP Cyg4178.85 QC,S forbidden FeI?4181.93 QD,S4184.07 QD,S4185.15 QD,S S4191.09 QD,S S4198.30 FeI 1.91E-01 2.399 QC > S P Cyg4200.7 QC > S nebular?4215.52 Sr II 7.16E-01 0.0 QC,S4216.18 FeI 4.41E-04 0.0 QC,S ec shifted blueward.4226.73 Ca I 1.75E+00 0.0 QC,S ecs shifted to the red4227.43 26.00 1.84E+00 3.332 QC,S P Cyg4233.15 QC,S4237.27 QC,S4254.35 Cr I 8.13E-01 0.000 QC,S4259.31 V I 6.76E-03 0.017 QC,S P Cyg4266.34 QC,S4272.23 QC,S4274.81 Cr I 6.03E-01 0.0 QC,S ec4277.538 TiII 1.816E-01 4.969 QC,S P Cyg4289.73 Cr I 4.27E-01 0.0 QC,S ec5889.95 Na I 1.28E+00 0.0 QC,S ecs5895.92 Na II 6.40E-01 0.0 QC,S ecs


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