Context. The chemical composition of the Sun is a fundamental yardstick in astronomy relative to which essentially all cosmic objects are referenced. As such having accurate knowledge of the solar elemental abundances is crucial for an extremely broad range of topics. BR  Aims We reassess the solar abundances of all 83 longlived elements using highly realistic solar modelling and stateoftheart spectroscopic analysis techniques coupled with the best available atomic data and observations. BR  Methods The basis for our solar spectroscopic analysis is a threedimensional 3D radiativehydrodynamical model of the solar surface convection and atmosphere which reproduces the full arsenal of key observational diagnostics. New complete and comprehensive 3D spectral line formation calculations taking into account of departures from local thermodynamic equilibrium nonLTE are presented for Na Mg K Ca and Fe using comprehensive model atoms with reliable radiative and collisional data. Our newly derived abundances for C N and O are based on a 3D nonLTE analysis of permitted and forbidden atomic lines as well as 3D LTE calculations for a total of 879 molecular transitions of CH CSUB2SUB CO NH CN and OH. Previous 3Dbased calculations for another 50 elements are reevaluated based on updated atomic data a stringent selection of lines improved consideration of blends and new nonLTE calculations available in the literature. For elements where spectroscopic determinations of the quiet Sun are not possible the recommended solar abundances are revisited based on complementary methods including helioseismology He solar wind data from the Genesis sample return mission noble gases sunspot observations four elements and measurements of the most primitive meteorites 15 elements. BR  Results Our new improved analysis confirms the relatively low solar abundances of C N and O obtained in our previous 3Dbased studies log SUBCSUB  8.46  0.04 log SUBNSUB  7.83  0.07 and log SUBOSUB  8.69  0.04. Excellent agreement between all available atomic and molecular indicators is achieved for C and O but for N the atomic lines imply a lower abundance than for the molecular transitions for unknown reasons. The revised solar abundances for the other elements also typically agree well with our previously recommended values with only Li F Ne Mg Cl Kr Rb Rh Ba W Ir and Pb differing by more than 0.05 dex. The hereadvocated presentday photospheric metal mass fraction is only slightly higher than our previous value mainly due to the revised Ne abundance from Genesis solar wind measurements XSUBsurfaceSUB  0.7438  0.0054 YSUBsurfaceSUB  0.2423  0.0054 ZSUBsurfaceSUB  0.0139  0.0006 and ZSUBsurfaceSUBXSUBsurfaceSUB  0.0187  0.0009. Overall the solar abundances agree well with those of CI chondritic meteorites but we identify a correlation with condensation temperature such that moderately volatile elements are enhanced by 0.04 dex in the CI chondrites and refractory elements possibly depleted by 0.02 dex conflicting with conventional wisdom of the past halfcentury. Instead the solar chemical composition more closely resembles that of the finegrained matrix of CM chondrites with the expected exception of the highly volatile elements. BR  Conclusions Updated presentday solar photospheric and protosolar abundances are presented for 83 elements including for all longlived isotopes. The socalled solar modelling problem  a persistent discrepancy between helioseismology and solar interior models constructed with a low solar metallicity similar to that advocated here  remains intact with our revised solar abundances suggesting shortcomings with the computed opacities andor treatment of mixing below the convection zone in existing standard solar models. The uncovered trend between the solar and CI chondritic abundances with condensation temperature is not yet understood but is likely imprinted by planet formation especially since a similar trend of opposite sign is observed between the Sun and solar twins.
['Sun: abundances', 'Sun: photosphere', 'Sun: atmosphere', 'line: formation', 'meteorites', 'meteors', 'meteoroids', 'Sun: helioseismology', 'Astrophysics - Solar and Stellar Astrophysics', 'Astrophysics - Earth and Planetary Astrophysics']
The chemical makeup of the Sun A 2020 vision
{'The chemical make-up of the Sun: A 2020 vision': 'M. Asplund ? , A. M. Amarsi 1 , and N. Grevesse 2 ; 3  \n- 1 Theoretical Astrophysics, Department of Physics and Astronomy, Uppsala University, Box 516, SE-751 20 Uppsala, Sweden\n- 2 Centre Spatial de Liège, Université de Liége, avenue Pré Aily, 4031 Angleur-Liège, Belgium\n- 3 Space Sciences, Technologies and Astrophysics Research (STAR) Institute, Université de Liège, Allée du 6 août, 17, B5C, 4000 Liège, Belgium  \nReceived: January 29, 2021; accepted: May 3, 2021', 'ABSTRACT': 'Context. The chemical composition of the Sun is a fundamental yardstick in astronomy, relative to which essentially all cosmic objects are referenced. As such, having accurate knowledge of the solar elemental abundances is crucial for an extremely broad range of topics.  \nAims. We reassess the solar abundances of all 83 long-lived elements, using highly realistic solar modelling and state-of-the-art spectroscopic analysis techniques coupled with the best available atomic data and observations.  \nMethods. The basis for our solar spectroscopic analysis is a 3D radiative-hydrodynamical model of the solar surface convection and atmosphere, which reproduces the full arsenal of key observational diagnostics. New complete and comprehensive three-dimensional (3D) spectral line formation calculations taking into account of departures from local thermodynamic equilibrium (non-LTE) are presented for Na, Mg, K, Ca, and Fe using comprehensive model atoms with reliable radiative and collisional data. Our newly derived abundances for C, N, and O are based on a 3D non-LTE analysis of permitted and forbidden atomic lines as well as 3D LTE calculations for a total of 879 molecular transitions of CH, C2, CO, NH, CN, and OH. Previous 3D-based calculations for another 50 elements are re-evaluated based on updated atomic data, a stringent selection of lines, improved consideration of blends, and new non-LTE calculations available in the literature. For elements where spectroscopic determinations of the quiet Sun are not possible, the recommended solar abundances are revisited based on complementary methods, including helioseismology (He), solar wind data from the Genesis sample return mission (noble gases), sunspot observations (four elements), and measurements of the most primitive meteorites (15 elements).  \nResults. Our new improved analysis confirms the relatively low solar abundances of C, N, and O obtained in our previous 3Dbased studies: log GLYPH<15> C = 8 : 46 GLYPH<6> 0 : 04, log GLYPH<15> N = 7 : 83 GLYPH<6> 0 : 07, and log GLYPH<15> O = 8 : 69 GLYPH<6> 0 : 04. Excellent agreement between all available atomic and molecular indicators is achieved for C and O, but for N the atomic lines imply a lower abundance than for the molecular transitions for unknown reasons. The revised solar abundances for the other elements also typically agree well with our previously recommended values, with only Li, F, Ne, Mg, Cl, Kr, Rb, Rh, Ba, W, Ir, and Pb di GLYPH<11> ering by more than 0.05 dex. The here-advocated present-day photospheric metal mass fraction is only slightly higher than our previous value, mainly due to the revised Ne abundance from Genesis solar wind measurements: X surface = 0 : 7438 GLYPH<6> 0 : 0054, Y surface = 0 : 2423 GLYPH<6> 0 : 0054, Z surface = 0 : 0139 GLYPH<6> 0 : 0006, and Z surface = X surface = 0 : 0187 GLYPH<6> 0 : 0009. Overall, the solar abundances agree well with those of CI chondritic meteorites, but we identify a correlation with condensation temperature such that moderately volatile elements are enhanced by GLYPH<25> 0 : 04 dex in the CI chondrites and refractory elements possibly depleted by GLYPH<25> 0 : 02 dex, conflicting with conventional wisdom of the past half-century. Instead, the solar chemical composition more closely resembles that of the fine-grained matrix of CM chondrites with the expected exception of the highly volatile elements.  \nConclusions. Updated present-day solar photospheric and proto-solar abundances are presented for 83 elements, including for all long-lived isotopes. The so-called solar modelling problem - a persistent discrepancy between helioseismology and solar interior models constructed with a low solar metallicity similar to that advocated here - remains intact with our revised solar abundances, suggesting shortcomings with the computed opacities and / or treatment of mixing below the convection zone in existing standard solar models. The uncovered trend between the solar and CI chondritic abundances with condensation temperature is not yet understood but is likely imprinted by planet formation, especially since a similar trend of opposite sign is observed between the Sun and solar twins.  \nKey words. Sun: abundances - Sun: photosphere - Sun: atmosphere - line: formation - Meteorites, meteors, meteoroids - Sun: helioseismology', '1. Introduction': "The pursuit of detailed knowledge of the elemental abundances in the Solar System has played a critical role in the development of modern astronomy. Since the chemical make-up of the planets is not expected to precisely reflect the composition of the nebula from which the Sun and planets formed, there are two methods to determine the proto-solar abundances, both of which were  \ninitiated a century ago: solar spectroscopy (Payne 1925; Russell 1929) and mass spectroscopy of primitive meteorites (Goldschmidt 1922, 1938). The two techniques each have their advantages and are thus highly complementary. Indeed a common approach is to combine the meteoritic and photospheric measurements to estimate the Solar System's composition, starting from the early influential works by Unsöld (1948) and Suess & Urey (1956), who provided a large amount of abundance data for the seminal papers on the origin of the elements by Burbidge et al.  \n(1957) and Cameron (1957). More recently, the widely used literature compilations by Anders & Grevesse (1989), Grevesse & Noels (1993), Grevesse & Sauval (1998), Lodders (2003, 2019), Lodders et al. (2009), and Palme et al. (2014) adopted this approach with a preference for the meteoritic values for most elements on account of their typically higher precision 1 . Due to severe depletion of the highly volatile elements - H, He, C, N, O, Ne, Ar, Xe, and Kr - meteorites cannot, however, be relied upon to measure the bulk Solar System composition, for which the Sun is indispensable. There is also no guarantee that the abundances of other elements in the most primitive meteorites, the CI chondrites, truly reflect those of the Sun; indeed, this long-held assumption is not correct, as we show here.  \nThe photospheric abundances are certainly not without shortcomings themselves. While often referred to as observed elemental abundances, the chemical composition of the Sun cannot be directly measured from the observed solar spectrum. Inferring the chemical make-up of a star requires a realistic model for the stellar atmosphere and of the transfer of radiation coupled with all atomic processes that shape the predicted emergent radiation to be compared with the observed spectrum. By necessity, the derived solar (and stellar) abundances are model dependent and therefore are not more reliable than the models from which they are inferred. Traditionally, quantitative solar spectroscopy has been done by means of one-dimensional (1D) model atmospheres: either theoretical models in radiative and convective equilibrium (e.g. Kurucz 1993; Gustafsson et al. 2008) or semi-empirical models in which the atmospheric temperature stratification is deduced from observations of the continuum centre-to-limb variation and of line depths, for example (e.g. Holweger & Müller 1974; Allende Prieto et al. 2001a; Fontenla et al. 2006). Neither type of 1D model attempts to describe the temporal variations or horizontal inhomogeneities in the atmosphere, and as such they include at best only a rudimentary description of convective energy transport (e.g. mixing length theory, Böhm-Vitense 1958).  \nA huge body of influential and careful work on determining the solar chemical composition has been carried out over several decades using such 1D solar atmosphere models (e.g. Goldberg et al. 1960; Lambert 1968, 1978; Grevesse & Sauval 1973, 1999; Ross & Aller 1976; Holweger et al. 1991; Blackwell et al. 1995). This has often been done in close collaboration with atomic physicists measuring or computing transition probabilities and other necessary input data, which is a highly successful synergy (e.g. Garz et al. 1969; Andersen et al. 1975; Biemont et al. 1981, 1991; Hannaford et al. 1992; Barklem et al. 2000; Lawler et al. 2001, 2013; Johansson et al. 2003; Ljung et al. 2006; Blackwell-Whitehead & Bergemann 2007; Den Hartog et al. 2019). These abundance analyses have mostly been within the framework of local thermodynamic equilibrium (LTE), thus assuming the Boltzmann, Saha, and Guldberg-Waage distributions for the atomic and molecular level populations and the Planck function for the radiation source function (e.g. Rutten 2003). While LTE is for many spectral lines that form in the solar photosphere a reasonable approximation, its validity cannot be taken for granted. For some elements, departures from LTE for the line formation in 1D model atmospheres (here: 1D non-LTE) and their impact on the inferred solar abundances  \nhave been explored (e.g. Steenbock & Holweger 1984; Kiselman 1991; Takeda 1994; Korn et al. 2003; Mashonkina et al. 2011; Lind et al. 2011; Bergemann et al. 2012; Shi et al. 2014). Unfortunately, such statistical equilibrium calculations have still not been performed for the majority of the elements, rendering their solar abundances prone to systematic errors (Asplund 2005; Bergemann et al. 2012; Nissen & Gustafsson 2018).  \nWith increased computational power, it has become possible to perform 3D (magneto-)hydrodynamical simulations of the outer convective and atmospheric layers in the Sun and other stars (e.g. Nordlund & Dravins 1990; Freytag et al. 1996; Ludwig et al. 1999; Asplund et al. 1999; Vögler et al. 2005; Collet et al. 2007; Muthsam et al. 2010; Magic et al. 2013; Trampedach et al. 2013; Beeck et al. 2013). In such modelling, the equations for the conservation of mass, momentum, and energy are solved together with the radiative transfer equation in a 3D simulation box covering a small but representative volume of the star. Besides shedding crucial light on, for example, the nature of convective energy transport (e.g. Stein & Nordlund 1998), the excitation and damping of solar and stellar oscillations (e.g. Sonoi et al. 2015; Zhou et al. 2019), and the dynamo-generation of stellar magnetic fields (e.g. Vögler & Schüssler 2007; Rempel 2014), such 3D simulations are eminently suited as model atmospheres for the prediction of the emergent radiation spectrum and thus for solar abundance analyses (e.g. Kiselman & Nordlund 1995; Asplund et al. 2000c, 2005b; Allende Prieto et al. 2001b; Holweger 2001; Ca GLYPH<11> au et al. 2008, 2011). Because convection emerges naturally in a self-consistent manner, there is no need to invoke the mixing length theory. These 3D hydrodynamical models are parameter-free in the sense that there is no tuning of any parameters involved to improve agreement with observations, yet their granulation properties, spectral line shapes, continuum centre-to-limb variations, and spectral energy distributions agree extremely well with observations; the choice of, for example, numerical resolution and sub-grid treatment does not impact the predictions (e.g. Asplund et al. 2000a; Nordlund et al. 2009; Pereira et al. 2013). Importantly, this is achieved without invoking any micro- and macro-turbulence parameters since the convective motions and stellar oscillations introduce the necessary Doppler shifts together with the (anti-)correlations in temperature and velocities in granules and intergranular lanes to explain the line strengths, widths, shifts and asymmetries (Asplund et al. 2000b). Clearly, the most recent 3D hydrodynamical solar models, such as the one employed herein, are highly realistic and far superior to any 1D model.  \nToday it is also feasible to carry out 3D non-LTE line formation calculations with comprehensive model atoms involving the relevant radiative and collisional processes for solar / stellar abundance purposes (e.g. Asplund et al. 2003, 2004; Cayrel et al. 2007; Lind et al. 2013; Ste GLYPH<11> en et al. 2015; Amarsi et al. 2016a, 2019; Bergemann et al. 2019). Such statistical equilibrium computations, while computationally demanding 2 , make optimal use of the highly realistic 3D hydrodynamical stellar surface convection models and self-consistently describe departures from the LTE within the trace-element assumption. The non-LTE e GLYPH<11> ects  \ncan typically be expected to be exacerbated in 3D compared to 1D modelling (Asplund 2005). The field has benefited tremendously from the increased availability of accurate atomic data for the vast number of radiative and collisional cross-sections required, including the previously highly uncertain inelastic collisions with electrons and with neutral hydrogen, which are of paramount importance in the atmospheres of late-type stars (e.g. Barklem 2016a). It is truly remarkable that all free parameters - mixing length parameters, microturbulence, macroturbulence, Unsöld enhancement factor for pressure broadening, Drawin scaling factor for H collisions, etc - which have hampered quantitative stellar spectroscopy for decades, are now finally obsolete. Together with vastly improved stellar modelling and transition probabilities, this has drastically improved the accuracy and precision of the inferred stellar chemical compositions.  \nAsplund et al. (2009) presented the first fully 3D-based analysis of the solar photospheric elemental abundances with nonLTE calculations (usually based on 1D model atmospheres) for the most important elements; earlier 3D determinations for many elements were reported in for example Asplund et al. (2000c, 2004, 2005a,b). It was also the first time all spectroscopically available elements were analysed in a homogeneous manner and with the systematic uncertainties quantified in detail (see also Scott et al. 2015a,b and Grevesse et al. 2015 for further details).  \nHere, we discuss further improvements to our solar analysis, including new 3D non-LTE calculations for many more elements, better atomic data, and refined line selection, which have led to a revision of the abundances of numerous elements. We argue that our study represents a state-of-the-art determination of the solar chemical composition, yielding the most reliable results available today. Complemented with precise laboratory measurements of the compositions of the most pristine meteorites (e.g. Palme et al. 2014), the new solar abundances enable us to uncover subtle abundance di GLYPH<11> erences with condensation temperatures of the elements between the Sun and CI chondrites, which reflect the complicated processes of star and planet formation. For convenience, we also provide estimates for the protosolar (isotopic) abundances where the e GLYPH<11> ects of atomic di GLYPH<11> usion, nuclear burning, and radioactive decay have been taken into account, as well as the present-day and proto-solar mass fractions of H, He, and heavy elements.", '2. Present-day photospheric solar abundances': 'The recommended solar photospheric elemental abundances presented in Table 2 and shown in Fig. 1 represent, in our opinion, the most reliable measurements available today, obtained with state-of-the-art analysis techniques and the best atomic and molecular data. All abundances have well-quantified and justified uncertainties and supersede those given in our previous summaries of the solar chemical composition (e.g. Asplund et al. 2005a, 2009). Table 1 intercompares the main ingredients of this study and our previous 3D-based solar abundance analyses in terms of 3D solar modelling and spectral line formation, showing that most recent progress has concerned non-LTE calculations, as described further below.  \nThe abundances are expressed on the traditional astronomical logarithmic abundance scale. Hydrogen is the natural reference element for solar (and stellar) spectroscopy, both because it is the most abundant element and because it provides the continuous opacity in the optical and infrared through the negative hydrogen ion H GLYPH<0> . The normalisation of the elemental number  \nFig. 1. Present-day solar photospheric logarithmic abundances (by number) as a function of atomic number, with H defined as log GLYPH<15> H GLYPH<17> 12 : 00. For elements for which no photospheric determination has been possible the CI chondritic abundance is shown, corrected for the identified correlation with condensation temperature (Sect. 3.3). H and He are the most abundant elements on account of being produced in the Big Bang with some contribution to He from stellar H-burning. Prominent peaks occur around O, Fe, Ba, and Pb, while elements with an even atomic number have higher abundances than neighbouring odd elements (so-called odd-even e GLYPH<11> ect) as a result of stellar evolution and nuclear physics (minimum nuclear binding energy and nuclei magic numbers). Li, F, and Sc have particularly low abundances relative to nearby elements on account of them being odd elements, having relatively low binding energy, and not being part of the main nucleosynthesis production channels in stars. We note that the unstable elements Tc ( Z = 43) and Pm ( Z = 61) are not shown.  \n<!-- image -->  \nAtomic number  \ndensity N X for an element X is defined as log GLYPH<15> X GLYPH<17> log ( N X = N H) + 12 : 00, for historical reasons, hence log GLYPH<15> H GLYPH<17> 12 : 00 : 3  \nAfter first summarising some central ingredients in our analysis, in the following sub-sections we present a detailed discussion of the elements for which solar photospheric abundances have been newly analysed, compared with our previous spectroscopic studies of elements heavier than Ne (Scott et al. 2015a,b; Grevesse et al. 2015). In some cases, we simply updated their results in light of new atomic data (in particular, transition probabilities) or improved non-LTE abundance corrections, or we made a di GLYPH<11> erent choice in terms of the line list or weighting of the individual lines. For Li, C, N, O, Na, Mg, Al, Si, K, Ca, Fe, and Ba, our results are based on new 3D non-LTE calculations. We also discuss the elements that are not accessible by spectroscopy of the quiet Sun for which other methods to infer the solar abundances are required.', '2.1. General considerations of the analysis': 'Model atmospheres: As a default solar model atmosphere, we employed the same 3D hydrodynamic simulation of the solar surface convection computed with the Stagger code (Nordlund & Galsgaard 1995; Stein & Nordlund 1998; Magic et al. 2013),  \nTable 1. Summary of main ingredients of our 3D-based solar abundance analyses.  \nas was done in our recent studies of 3D non-LTE line formation of C, N, and O in the Sun (Amarsi et al. 2018a, 2019, 2020). This is an updated version of the Stagger solar model used in Asplund et al. (2009), Scott et al. (2015a,b), and Grevesse et al. (2015), but the di GLYPH<11> erences in photospheric structure are minor. Detailed 3D radiative transfer with line-blanketing through the opacity binning technique (Nordlund 1982; Collet et al. 2018) is included under the assumption of LTE to compute the necessary radiative heating and cooling rates in the photosphere. The resulting temperature stratification is highly realistic, as is evident from comparisons with the observed continuum centreto-limb variation (Amarsi et al. 2018a); this was also the case with the 3D model used in Asplund et al. (2009). The predictions of our 3D solar models are in excellent agreement with all key observational diagnostics (e.g. Pereira et al. 2013). The resulting e GLYPH<11> ective temperature of the 24 hr simulation sequence is 5773 GLYPH<6> 16 K (standard deviation), indistinguishable from the solar value of 5772 K (Prša et al. 2016). The 3D Stagger models employ an updated version of the so-called MHD (named after the authors) equation of state (Mihalas et al. 1988; Trampedach et al. 2013), continuous and line opacities from the MARCS package (Gustafsson et al. 2008; Hayek et al. 2010), and the solar chemical composition of Asplund et al. (2009). No magnetic fields were included. From the 240 3 simulation grid covering 6 Mm horizontally and 3 Mm vertically, a resampled 3D grid covering the photosphere (1 Mm vertically) was interpolated before the spectral line formation calculations for the abundance determinations to improve the numerical accuracy and to save computational time without sacrificing accuracy; the grid resolutions of both the original and interpolated solar models are su GLYPH<14> cient to obtain the correct spectral line broadening from convective motions and oscillations (Asplund et al. 2000a,b; Pereira et al. 2013).  \nTo quantify various systematic uncertainties, we also made use of several 1D model atmospheres: the temporally and spatially averaged 3D model (referred to here as < 3D > ); a 1D hy-  \ndrostatic ATMO model with mixing length theory convection and identical equation of state and opacities to the 3D model (Magic et al. 2013, Appendix A); a 1D hydrostatic model computed with the widely used MARCS code (Gustafsson et al. 2008); and the 1D semi-empirical Holweger & Müller (1974) model, which has traditionally been the model of choice for solar abundance analyses for the past half-century.  \nSpectral line formation: For key atoms and ions, the results are based on 3D non-LTE calculations, using the MPIparallelised Balder code (Amarsi et al. 2018b), our extensively modified version of Multi3D (Leenaarts & Carlsson 2009). For Li i , C i , N i , Al i , and Si i / Si ii , these results were adopted from recently published studies (Wang et al. 2021; Amarsi et al. 2019, 2020; Nordlander & Lind 2017; Amarsi & Asplund 2017), while for O i , Na i , Mg i / Mg ii , K i , Ca i / Ca ii , and Fe i / Fe ii , they are based on new calculations. We also adopted 3D non-LTE literature results for Ba ii , which are based on Multi3D calculations (Gallagher et al. 2020).  \nBalder iteratively solves the 3D (or 1D) radiative transfer equation coupled with the equations for statistical equilibrium (i.e. the atomic level populations fulfil @ n i =@ t = 0) under the assumption that the departures from LTE for an element do not feed back onto the atmospheric structure, nor do they impact the line formation of other elements. Balder employs a comprehensive equation of state, background continuum, and line opacities calculated with the Blue code as described in Amarsi et al. (2016b), with the present work adopting the occupation probability formalism to better model line opacities near the series limits (Hubeny et al. 1994). The 3D non-LTE radiative transfer is solved iteratively for typically 5-10 snapshots of the 3D solar model in 26 directions using short characteristics (Ibgui et al. 2013). The number of frequencies for which the radiative transfer equation needs to be computed to adequately resolve the bound-free and bound-bound transitions varies among elements, but for our most complex model atoms (Ca and Fe) it amounts to 20000 and 45000, respectively. After convergence, a final formal  \nsolution is carried out for the emergent vertical intensity and for the spectral lines of interest, using a long characteristics solver and higher frequency resolution for accuracy.  \nOther atomic and ionic lines, as well as all molecular lines, are based on 3D LTE spectrum synthesis using the SCATE code (Hayek et al. 2010), for the most part as presented in our earlier papers (Scott et al. 2015a,b; Grevesse et al. 2015). New 3D LTE calculations were adopted for molecular species of C, N, and O. When no 3D non-LTE calculations had been performed, we relied on our own or published 1D non-LTE abundance corrections, if possible; this procedure is not fully consistent but is expected to be more realistic than the LTE results. Many elements, in particular those beyond the Fe peak, unfortunately still lack non-LTE studies, requiring us to adopt the 3D LTE results. It is emphasised that no micro- and macro-turbulence enter the 3D calculations (Asplund et al. 2000b); for the comparison 1D spectrum syntheses, we assumed a microturbulence of 1 km s GLYPH<0> 1 (e.g. Holweger et al. 1978; Ste GLYPH<11> en et al. 2013).  \nAtomic data: The computation of the theoretical solar spectrum hinges critically on the availability of accurate atomic data. We have scoured and carefully assessed the literature and databases for the most up-to-date and most reliable input data required for our solar modelling: transition probabilities, hyperfine splitting, spectral line broadening, partition functions, excitation, ionisation, and dissociation energies, etc (e.g. Piskunov et al. 1995; Barklem & Collet 2016; Heiter et al. 2021). In most cases, experimental data have been selected over theoretical data, with the NIST Atomic Spectra Database 4 (Kramida 2019) providing much of the quantitative assessment of the uncertainties in gf -values. For the solar photosphere, the continuous opacities are well determined, except possibly in the UV, and hence they are not a major concern. For the non-LTE calculations, we also made extensive use of the bound-bound and bound-free radiative data from the Opacity Project, the Iron Project, and other dedicated atomic calculations (e.g. Seaton 1987; Hummer et al. 1993; Badnell et al. 2005), as well as theoretical cross-sections for collisions with electrons and with neutral hydrogen (e.g. Bray & Stelbovics 1992, 1995; Zatsarinny 2006; Zatsarinny & Bartschat 2013; Barklem 2016b; Barklem et al. 2017; Belyaev & Yakovleva 2017). Fortunately the situation today in terms of the necessary atomic and molecular data for solar / stellar spectroscopy is significantly improved compared to just a decade ago thanks to the tireless contributions from a relatively small group of atomic physicists to whom we are greatly indebted.  \nObservations: To derive the solar elemental abundances, the predicted solar spectrum was compared to the observed highresolution disc-centre intensity spectrum of the quiet Sun with a preference for the Liège (named after its location of production, not observations which were done at Jungfraujoch) visual solar atlas (Delbouille et al. 1973) due to its low degree of contamination of telluric absorption lines. We also made use of other solar atlases for the disc-centre intensity (Neckel 1999; Kurucz 2005), the disc-integrated flux (Kurucz et al. 1984; Kurucz 2006; Reiners et al. 2016), as well as observations in the infrared (Delbouille et al. 1981; Farmer & Norton 1989; Abrams et al. 1996; Hase et al. 2010) as needed and for evaluation purposes; disccentre intensity spectra are in general preferable to flux to disentangle the contribution of blends. In some cases, we also employ the high-resolution optical centre-to-limb observations of Pereira et al. (2009) and Stenflo (2015) to verify the correctness of our line formation calculations. All of these observed solar  \nspectra have extremely high resolving power (typically equivalent of R GLYPH<17> GLYPH<21>= GLYPH<1> GLYPH<21> > 200 ; 000) and signal-to-noise (mostly S = N > 1000). We either perform a line profile fitting using standard GLYPH<31> 2 -techniques or match the observed and computed equivalent widths, depending on the line and possible presence of blends. We allowed our own local continuum placement around the spectral lines in question, as appropriate. Doerr et al. (2016) and Reiners et al. (2016) examined the di GLYPH<11> erences between some of these atlases and overall find excellent agreement; our results are not significantly a GLYPH<11> ected by the particular choice of observations except for a few individual lines.  \nUncertainties: Homogenously quantifying the uncertainties of elemental abundances, including systematic errors is almost as important as having accurate determinations of them. Concerning elements for which multiple spectral lines are available for the abundance analysis, we used the weighted standard error of the mean as the statistical uncertainty, the weights being our assessment of the quality of the lines. In cases when only one or two lines can be used, we estimated the statistical error from the goodness of the profile fitting or equivalent width measurement uncertainty.  \nFollowing Asplund et al. (2009), we devoted much e GLYPH<11> ort to carefully investigating and evaluating possible systematic errors in the analysis, in particular related to the 3D model atmosphere and 3D non-LTE line formation. To quantify the errors from shortcomings in the photosphere temperature structure we take half the di GLYPH<11> erences between the derived abundances from the < 3D > model and the Holweger & Müller (1974) model; as noted above, we know from the centre-to-limb variation that the < 3D > model is a much closer match to the real Sun (Asplund et al. 2009; Pereira et al. 2013) with newer semi-empirical solar model atmospheres (Grevesse & Sauval 1998; Allende Prieto et al. 2001a) having 50 GLYPH<0> 100 K lower temperatures than the Holweger & Müller (1974) model in the line-forming region, which is very similar to our 3D model. The impact of atmospheric inhomogeneities is evaluated from taking half the di GLYPH<11> erence between the 3D and < 3D > results, which is justified by the close resemblance between the predicted and observed solar granulation patterns. In general, the systematic uncertainty due to non-LTE line formation is estimated as half of the predicted non-LTE abundance correction (3D non-LTE GLYPH<0> 3D LTE, or 1D non-LTE GLYPH<0> 1D LTE) with a minimum error of 0.03 dex when no non-LTE calculations are available; an exception is made for O for which the non-LTE e GLYPH<11> ects can be accurately quantified using, for example, the centre-to-limb variation of the O i 777 nm triplet and therefore the uncertainties would otherwise be significantly overestimated. These three systematic errors are combined in quadrature with the statistical error to compute the total abundance error.  \nFor some elements, we tried to factor in uncertainties in the input atomic data but in general this has not been attempted due to those often being poorly known and the huge amount of atomic data required for non-LTE modelling. While impossible to know without even more sophisticated solar modelling, we expect our estimated uncertainties to be realistic and possibly conservative overall, although we caution that in a few cases, especially in the absence of any non-LTE calculations, the errors are likely underestimated; in cases where they were particularly striking, we artificially inflated the computed errors somewhat, as noted.  \nTable 2. Present-day solar photospheric abundances together with those in CI chondrites (Palme et al. 2014), renormalised to the solar Si abundance (Sect. 3). Elements that are commented on are those for which the adopted proto-solar abundances (Sect. 5) are not based on spectroscopy of the solar photosphere; we caution that the di GLYPH<11> erences between solar and CI abundances show a trend with condensation temperature (Sect. 3.3), which should be taken into account when estimating the Solar System composition.', '2.2. Light elements: Lithium to beryllium': 'Li, Be, and B were not analysed in our previous series of papers devoted to the solar chemical composition (Scott et al. 2015a,b; Grevesse et al. 2015) and are therefore discussed explicitly here.  \nLithium (Z = 3): In practise, the solar Li abundance can only be determined from the very weak Li i 670.8 nm resonance line, an undertaking that is further complicated by the presence of several blending lines of CN, Fe, and other species. We adopted the recently revised solar Li abundance from Wang et al. (2021), who presented new 3D non-LTE spectral line formation calculations for a wide range of stellar parameters, including for the Sun. Importantly, their more complete treatment of background line opacities in the UV compared with previous studies (e.g.  \nCarlsson et al. 1994; Asplund et al. 2003; Lind et al. 2009, 2013; Harutyunyan et al. 2018; Mott et al. 2020) significantly changes the statistical equilibrium balance of Li i , leading in general to a strengthening of the 670.8 nm line. Using their new 3D non-LTE calculations and allowing for blends, Wang et al. (2021) derived a solar Li abundance of log GLYPH<15> Li = 0 : 96 GLYPH<6> 0 : 03 (statistical) GLYPH<6> 0 : 05 (systematic) and a 0.09 dex (23%) reduction compared to the advocated value in Asplund et al. (2009).  \nBeryllium (Z = 4): Following Chmielewski et al. (1975), it was widely believed that Be was depleted by a factor of two in the solar photosphere compared with CI chondrites. This conclusion was first challenged by Balachandran & Bell (1998) who demonstrated that missing UV opacity (e.g. Bell et al. 2001;  \nShort & Hauschildt 2009) in the stellar models leads to an underestimated solar Be abundance when relying on the Be ii 313 nm resonance lines. By calibrating the missing opacity in this wavelength region by enforcing the same O abundance to be estimated from nearby UV OH lines with OH lines in the IR, they found a solar photospheric Be abundance consistent with the meteoritic value. Their conclusion was corroborated in a 3D LTE analysis by Asplund (2004), which we adopted (log GLYPH<15> Be = 1 : 38 GLYPH<6> 0 : 09; statistical + systematic), suggesting no solar Be depletion. We caution, however, that unaccounted-for blends in this crowded UVregion may have biased the inferred Be abundance upwards. Revisiting the topic with 3D non-LTE calculations and a careful assessment of blending lines and continuum placement together with a detailed evaluation of any missing UV opacity would be worthwhile to finally settle the issue, especially given the importance of accurately constraining the amount of mixing below the solar convection zone (Sect. 3.2).  \nBoron (Z = 5): The solar B abundance can in practice only be inferred from the B i 249.7 nm resonance line, which makes the result quite uncertain due to the heavy crowding and possible missing opacity in this wavelength region. The most recent spectroscopic analysis is that of Cunha & Smith (1999), who inferred log GLYPH<15> B = 2 : 70 GLYPH<6> 0 : 20 from a 1D LTE spectrum synthesis, which we have consequently adopted here as the present-day photospheric abundance. While the minority ionisation stage and low excitation potential may suggest that the 249.7 nm line should be susceptible to severe departures from LTE, this does not seem to be the case for the Sun due to a somewhat fortuitous cancellation of non-LTE e GLYPH<11> ects in 1D (Kiselman 1994; Tan et al. 2010). No 3D non-LTE study for B exists to our knowledge.', '2.3. Carbon, nitrogen, and oxygen': 'The fourth, sixth, and third most abundant elements in the Sun by number, C, N, and O are incredibly important to modern astronomy. They also play a central role in the outstanding solar modelling problem (e.g. Basu & Antia 2008; Zhang et al. 2019), wherein a highly significant discrepancy exists pertaining to the interior structure of the Sun as inferred precisely via helioseismic inversions and standard solar models (see Sect. 6). O is the dominant source of opacity at the most problematic region, namely the base of the convection zone, with C and N also providing a non-negligible contribution there. It has been noted that higher C, N, and (especially) O abundances by GLYPH<25> 0 : 15 dex would alleviate the discrepancy with the sound speed (e.g. Pinsonneault & Delahaye 2009), but this would not address all observational constraints (Buldgen et al. 2019b). Adding to the controversy, all three elements form volatile compounds and are thus depleted in meteorites; their abundances must be inferred through spectroscopy of the solar photosphere.  \nParallel to our work, the CO 5 BOLD group have led a commendable e GLYPH<11> ort to measure the solar C, N, and O abundances from atomic lines, similarly employing 3D hydrodynamic model atmospheres (Ca GLYPH<11> au et al. 2008, 2009, 2010, 2013, 2015; Steffen et al. 2015); they have not yet explored the molecular lines. These studies tend to find slightly higher C, N, and O abundances by 0 : 03 to 0 : 07 dex that usually agree within the uncertainties (Ca GLYPH<11> au et al. 2011). These di GLYPH<11> erences most likely do not originate in the 3D models or radiative transfer codes themselves, but rather can be traced to our more conservative choice of unblended lines, the adopted oscillator strengths, and the measured equivalent widths, as has been demonstrated for C (Amarsi et al. 2019) and N (Amarsi et al. 2020); moreover, unpublished cal-  \nculations have similarly indicated good agreement for the O i 777 nm triplet when adopting exactly the same input data.  \nAfter submission of our article, we became aware of new 3D non-LTE line formation calculations of the O i 777 nm triplet and [O i ] 630 nm forbidden line by Bergemann et al. (2021) using 3D solar models computed with the Stagger and Bifrost codes, the latter with and without a chromosphere. Their results agree perfectly with those presented here and in Asplund et al. (2009): log GLYPH<15> OI777 = 8 : 70 GLYPH<6> 0 : 05 and log GLYPH<15> [OI] 630 = 8 : 71 GLYPH<6> 0 : 03, further supporting a low solar O abundance. Reassuringly, they find only a very minor e GLYPH<11> ect ( GLYPH<25> 0 : 01 dex) of a chromospheric temperature rise on these lines.  \nAn extension to using a 1D semi-empirical model atmosphere has recently been presented by Cubas Armas et al. (2017, 2020), who derived a 3D solar model from the spectral inversion of three spatially resolved Fe i lines. They applied it to the [O i ] 630 : 0 nm line in the Sun to infer log GLYPH<15> O = 8 : 80 GLYPH<6> 0 : 03. While this is an interesting approach well worth pursuing further, we consider it premature to adopt their derived abundance until this empirical 3D solar model has been exposed to a variety of observational tests, which our theoretical 3D hydrodynamical models have successfully passed (Nordlund et al. 2009; Pereira et al. 2013). In particular, it is not clear whether their model has the appropriate T e GLYPH<11> since it is neither an imposed constraint of the spectral inversion of the normalised line profiles, nor is it tested for afterwards. It remains to be seen whether consistent O abundances would be obtained from the [O i ], O i , and OH lines.  \nTable 3 shows the results of our reanalysis of all the available indicators of C, N, and O using our most recent 3D model atmosphere and the most up-to-date atomic and molecular data. The atomic results are based on recent 3D non-LTE studies for C (Amarsi et al. 2019), N (Amarsi et al. 2020), and O (Amarsi et al. 2018a); in the latter case, the calculations were extended to include the other available atomic indicators (see also Amarsi 2016). The molecular results come from a new homogeneous analysis of 879 di GLYPH<11> erent lines, with the background C, N, and O abundances iterated to consistency (Amarsi et al. 2021). Compared to Asplund et al. (2009), the new recommended abundances are identical for N and O and only 0.03 dex higher for C.  \nCarbon (Z = 6): Of the three elements considered in this section, C has the largest number of available indicators (Table 3) and is arguably the best constrained. The [C i ] result is drawn from the 3D non-LTE analysis of the 872 : 7 nm line from Amarsi et al. (2019), although we note that this low-excitation line is well-reproduced in 3D LTE. The C i result is based on the 14 lines also analysed in this 3D non-LTE study; here, the result has increased by 0 : 03 dex in light of new gf -values from large-scale atomic structure calculations by Li et al. (2021). The molecular results are based on new 3D LTE calculations of 39 lines in the C2 Swan system, and CH lines which were divided into 51 fundamental rovibrational (d ν = 1) lines and seven electronic lines in the CH A-X system. CO lines were also considered by adopting the O abundance inferred from OH lines for a particular model atmosphere and including the error in this O abundance in the total error budget. The CO lines were divided into two groups of rovibrational transitions, including 28 belonging to fundamental (d ν = 1) bands and 52 to first overtone (d ν = 2) bands.  \nThe advocated solar C abundance is 8 : 46 GLYPH<6> 0 : 04. This was evaluated as the weighted mean and error of the [C i ], C i , C2, CH, and CO results given in Table 3; to avoid giving these species too much weight, the rovibrational and electronic results for CH were first combined into a weighted mean (8 : 46), as were  \nTable 3. C, N, and O abundances in the solar photosphere as inferred from the available indicators and from di GLYPH<11> erent model atmospheres, with the combined statistical and systematic uncertainties in the final column.  \nthe fundamental and first overtone results for CO (8 : 48). The final result is a 0 : 03 dex increase over those stated in Asplund et al. (2009). Nevertheless, the agreement between the di GLYPH<11> erent indicators found here is excellent, and noticeably improved over that earlier study. The reason for the new, slightly higher value can primarily be attributed to improvements to the atomic and molecular input data ([C i ] Froese Fischer 2006; C i Li et al. 2021; C2 Brooke et al. 2013; CH Masseron et al. 2014; CO Li et al. 2015), as well as to the employed molecular dissociation constants (Barklem & Collet 2016).  \nNitrogen (Z = 7): Compared to C, there are relatively few good indicators of the solar N abundance. There are no [N i ] lines visible in the solar spectrum, and only five N i lines were identified as being su GLYPH<14> ciently strong and unblended to be measured with confidence (Amarsi et al. 2020). The 3D non-LTE atomic result is combined with new 3D LTE results for molecular lines (Table 3). This includes NH rovibrational lines that were divided into two groups consisting of 13 pure rotational (d ν = 0) and 15 fundamental (d ν = 1) lines. CN lines were also considered by adopting weighted mean carbon abundances as inferred from C2, CH, and CO and for the di GLYPH<11> erent model atmospheres, and including the error in this C abundance in the total error budget. The CN lines were separated into two groups consisting of 59 electronic lines in the 0 GLYPH<0> 0 band, which typically has the best data, and 463 more lines with d ν GLYPH<21> 1 in various bands.  \nTable 3 shows a discrepancy of 0 : 13 dex between the atomic N i result (7 : 77), and the weighted mean (7 : 90) of the mean NH (7 : 91) and CN (7 : 88) results. We can only speculate on where the inconsistency arises. The N i oscillator strengths were taken from Tachiev & Froese Fischer (2002) and are expected to be reliable at the 0 : 03 dex level based on rankings from NIST (and verified with separate calculations as discussed in Sect. 2.1 of Amarsi et al. 2020). The oscillator strengths used here for NH and CN were taken from Brooke et al. (2014a,b, 2015), however it is di GLYPH<14> cult to imagine both sets of calculations having systematic o GLYPH<11> sets of 0 : 1 dex. Moreover, the dissociation energies of CN and NH are now well established (Barklem & Collet 2016) and  \nthus unlikely to be the reason behind the high molecular-based N abundance. We note that this discrepancy was already present in Asplund et al. (2009), albeit to a lesser degree, with the calculations there based on an unpublished set of older molecular data (A.J. Sauval, private communication).  \nThe advocated solar N abundance is 7 : 83 GLYPH<6> 0 : 07. Owing to the atomic versus molecular discrepancy, this was found by taking the unweighted mean of the atomic result together with the weighted mean of the molecular results. The uncertainty is given by the standard error in the mean (or equivalently, half the range between the two values). The value is unchanged from that of Asplund et al. (2009); however, the stipulated uncertainty is 0 : 02 dex larger.  \nOxygen (Z = 8): The solar O abundance can be inferred from a number of di GLYPH<11> erent atomic and molecular indicators (Table 3). For the atomic indicators, we carried out a spectral line fitting analysis of the [O i ] 557 : 7 nm, 630 : 0 nm, and 636 : 4 nm lines, and of the O i 615 : 8 nm, 777 nm, 844 : 7 nm, and 926 : 9 nm lines, using the 3D non-LTE model of Amarsi et al. (2018a). The fitting method follows that presented in Amarsi (2016, Chapter 4). Here, the strengths of the CN and C2 blending lines were fixed by adopting the line lists and abundances of the earlier C and N analyses and the strength of the blending Ni i 630 nm line calculated using the gf -value from Johansson et al. (2003) and the Ni abundance inferred here (log GLYPH<15> Ni = 6 : 20; Sect. 2.5). Ultimately, for [O i ] nearly all weight was given to the 630 : 0 nm line after considering the relative di GLYPH<14> culties in modelling the respective continua and blending contributions. Similarly for O i , most weight was given to the 777 nm triplet, because available centreto-limb observations (Pereira et al. 2009) allow us to correct for and constrain errors in the cross-sections for inelastic collisions with neutral hydrogen (following Amarsi et al. 2018a). The molecular results are based on new 3D LTE calculations of OH rovibrational lines, which were divided into three groups consisting of 84 pure rotational (d ν = 0), 50 fundamental (d ν = 1), and 15 first overtone (d ν = 2) lines, all giving highly consistent results.  \nThe advocated solar O abundance is 8 : 69 GLYPH<6> 0 : 04. Our new value was determined from the weighted mean and error of the [O i ], O i , and OH results given in Table 3, albeit first combining the three rovibrational OH results into a single weighted mean (8 : 70). An alternative value is obtained by taking the unweighted mean, either of these three species, or of all five indicators in Table 3. This would put less weight on the strong O i 777 nm triplet, but with both approaches the value is increased by less than 0 : 005 dex, since there is excellent agreement between the di GLYPH<11> erent indicators. Our recommended solar O abundance is unchanged from Asplund et al. (2009); however, the stipulated uncertainty is 0 : 01 dex smaller.  \nCombined with our revised C abundance, the solar C / Oratio is 0 : 59 GLYPH<6> 0 : 08, a value which has steadily increased over the last few decades (e.g. 0 : 43: Anders & Grevesse 1989; 0 : 48; Grevesse et al. 1991; 0 : 49: Grevesse & Sauval 1998; 0 : 50: Allende Prieto et al. 2001b, 2002; 0 : 54: Asplund et al. 2004, 2005b; 0 : 55: Asplund et al. 2009; Ca GLYPH<11> au et al. 2011; 0 : 56: Amarsi et al. 2018a, 2019). The Sun has a higher C / O ratio than most solar twins regardless of age (Bedell et al. 2018), which could be a result of the Sun having migrated radially outwards by 2 GLYPH<0> 3 kpc in the Milky Way disc since birth (Nieva & Przybilla 2012).', '2.4. Intermediate-mass elements: Sodium to calcium': "Sodium (Z = 11): The solar Na abundances presented in Asplund et al. (2009) and Scott et al. (2015b) were based on a 3D LTE analysis corrected for non-LTE e GLYPH<11> ects using 1D model atmospheres (Lind et al. 2011). Here, we present the first full 3D non-LTE calculations for Na in the Sun, employing the same 23-level Na model atom as in Lind et al. (2011), which importantly contains a full quantum mechanical description for inelastic Na i + H i collisions (Belyaev et al. 2010; Barklem et al. 2010). Experiments suggest these are su GLYPH<14> ciently realistic for non-LTE modelling (Barklem et al. 2021).  \nAsummary of the inferred abundances for the di GLYPH<11> erent lines using a variety of model atmospheres is provided in Table A.1. We employed the same five weak Na i lines as well as transition data as in Scott et al. (2015b); the equivalent widths were corrected for minor contributions from blending lines. The 3D nonLTE e GLYPH<11> ects are minor ( GLYPH<25> GLYPH<0> 0 : 03 dex) and very similar to those previously determined in 1D. Our recommended 3D non-LTE Na abundance, based on the unweighted mean of the five Na i lines, becomes log GLYPH<15> Na = 6 : 22 GLYPH<6> 0 : 01 (statistical) GLYPH<6> 0 : 03 (systematic), which is 0 : 01 dex larger than the value provided in Scott et al. (2015b).  \nMagnesium (Z = 12): Scott et al. (2015b) determined the solar photospheric Mg abundance based on 3D LTE line formation calculations coupled with non-LTE abundance corrections computed with di GLYPH<11> erent 1D model atmospheres. Their value, log GLYPH<15> Mg = 7 : 59 GLYPH<6> 0 : 04, was only 0.01 dex lower than the recommended abundance by Asplund et al. (2009). The Mg i and Mg ii lines yielded fully consistent abundances in non-LTE, but they were significantly less so in LTE, highlighting the need for statistical equilibrium calculations for Mg ii (non-LTE corrections are vanishingly small for Mg i for the Sun). Osorio et al. (2015) and Bergemann et al. (2017) carried out new 1D non-LTE calculations based on improved atomic data, but since their analyses were more focussed on other stars they have only very few lines in common with Scott et al. (2015b), most of their lines being much stronger in the Sun and less suitable for abundance purposes.  \nFor the first time, we carried out full 3D non-LTE calculations of Mg in the solar photosphere. Our new atomic model  \nconsists of 48 Mg i levels, 27 Mg ii levels, and the Mg iii ground state; and 482 bound-bound and 71 bound-free radiative transitions, drawing on new experimental and theoretical transition probabilities for Mg i (Pehlivan Rhodin et al. 2017b). We made use of realistic cross-sections for inelastic collisions with electrons from Barklem et al. (2017) for Mg i and from Liang et al. (2009) for Mg ii , and we used rate coe GLYPH<14> cients from Barklem et al. (2012) and Barklem (2016b) to describe inelastic collisions between Mg i and neutral hydrogen.  \nAsummary of the inferred abundances for the di GLYPH<11> erent lines using a variety of model atmospheres is provided in Table A.1. The unweighted mean of eight Mg i lines gives log GLYPH<15> Mg = 7 : 56 GLYPH<6> 0 : 01 (statistical) in 3D non-LTE. The systematic uncertainty as estimated from the 3D LTE, < 3D > , and HM models is less than 0 : 01 dex as these all give highly consistent results for our Mg i lines. In 3D LTE, our Mg i result is 0 : 04 dex smaller than that given in Scott et al. (2015b) with the new gf -values only having a minor e GLYPH<11> ect ( GLYPH<0> 0 : 01 dex). Instead, the di GLYPH<11> erence stems mainly from how pressure broadening was treated; here, we used ABO theory (Anstee & O'Mara 1995; Barklem & O'Mara 1997; Barklem et al. 1998), adopting new and extended tables (P. S. Barklem, private communication 5 ). Accordingly, we increased the systematic uncertainty in the Mg i result to 0 : 02 dex.  \nThe 3D non-LTE result from the unweighted mean of six lines of Mg ii is significantly smaller: log GLYPH<15> Mg = 7 : 52 GLYPH<6> 0 : 03 (statistical). The Mg ii results have a noticeably larger line-toline scatter, beyond the uncertainties in the adopted gf -values, which are all classified A or A + ( GLYPH<14> log gf GLYPH<20> GLYPH<6> 0 : 01 dex) by NIST (Kramida 2019). The systematic uncertainty as estimated from the 3D LTE, < 3D > , and HM models is 0 : 03 dex and is primarily driven by the large departures from LTE, in particular for three of the lines. Notably, we found that these lines are highly sensitive to the adopted recipe for inelastic collisions between Mg ii and neutral hydrogen. Our results here are based on the Drawin (1968) recipe as formulated by Lambert (1993), without any ad-hoc scale factors applied, for lack of a better approach. Since the 3D LTE result is log GLYPH<15> Mg = 7 : 58, more e GLYPH<14> cient collisions would help bring the Mg ii result into better agreement with Mg i . Another uncertainty concerns the choice of Mg abundance for calculating the equation of state and hence electron pressures. Here, this abundance was varied consistently with the line formation calculations. An alternative approach would fix the Mg abundance to the value used to construct the model atmosphere (log GLYPH<15> Mg = 7 : 60) and vary the gf -values instead. Test calculations on the < 3D > model suggest that this alternative approach could increase the overall Mg ii result by almost 0 : 02 dex; the electron pressure does not a GLYPH<11> ect Mg i to first order since it is the minority ionisation stage in the solar photosphere.  \nOur recommended 3D non-LTE Mg abundance is the average of Mg i and Mg ii , weighted by the total statistical and systematic errors: log GLYPH<15> Mg = 7 : 55 GLYPH<6> 0 : 03. A case could also be made for only adopting the Mg i results given the additional complications with Mg ii , but the final abundance would only be 0.01 dex higher.  \nAluminium (Z = 13): Nordlander & Lind (2017) recently performed full 3D non-LTE spectral line formation calculations for Al i lines in the Sun, based on their new model atom that includes realistic cross-sections for inelastic collisions with neutral hydrogen (Belyaev 2013). They employed the same 3D hydrodynamical Stagger model solar atmosphere and 3D non-LTE radiative transfer code Balder as we employed in this study. Impressively, their solar modelling also agrees well with the ob-  \nserved IR emission lines of Al i . They adopted the same lines and line data as in Scott et al. (2015b), except that they excluded the 1089.1 nm line due to telluric contamination, finding a solar photospheric Al abundance of log GLYPH<15> Al = 6 : 43 GLYPH<6> 0 : 03 (statistical + systematic). Here, we adopted the Nordlander & Lind (2017) solar abundance, which is identical to the 3D LTE results of Scott et al. (2015b); for the Sun, departures from LTE are in fact largely unimportant for these five Al i lines.  \nSilicon (Z = 14): Here, we adopted the solar Si abundance from our recent analysis (Amarsi & Asplund 2017), the first full 3D non-LTE line formation calculations performed for this element. The same atomic data and line weighting as in Scott et al. (2015b) were used for the nine Si i and one Si ii lines. The inferred photospheric Si abundance (log GLYPH<15> Si = 7 : 51 GLYPH<6> 0 : 03) is identical to the values found by Asplund et al. (2009) and Scott et al. (2015b); this value is used to normalise the meteoritic abundance scale to the photospheric one, see Sect. 3.1.  \nIt should be noted that Pehlivan Rhodin et al. (2021) recently computed new gf -values for a large number of Si i and Si ii lines, which are partly validated against experimental data. Had we adopted those for the seven Si lines in common with our study, the photospheric abundance would have increased to log GLYPH<15> Si = 7 : 53 GLYPH<6> 0 : 03 with a slightly reduced statistical scatter and a corresponding adjustment of the meteoritic abundances. We did, however, opt to retain the older experimental transition probabilities (Garz 1973), renormalised to the lifetimes of Becker et al. (1980) and O'Brian & Lawler (1991)), mainly since the new theoretical values imply a significant ionisation imbalance of 0.1 dex between Si i and Si ii , even in 3D non-LTE. The overall agreement with the meteoritic abundances further supports the Garz (1973) data (see Sect. 3.2). Further theoretical and experimental work on Si would be very welcome.  \nPhosphorus (Z = 15): We adopted the results from the 3D LTE analysis of Scott et al. (2015b): log GLYPH<15> P = 5 : 41 GLYPH<6> 0 : 03. To our knowledge, no non-LTE investigation of P exists in the literature, neither in 1D nor 3D. Due to the similarity between the weak P i and S i lines, the latter for which non-LTE investigations do exist (see below), comparably small non-LTE e GLYPH<11> ects could be expected for the two elements. The value is unchanged from Asplund et al. (2009).  \nSulphur (Z = 16): The solar S abundance log GLYPH<15> S = 7 : 12 GLYPH<6> 0 : 03 is taken from Scott et al. (2015b) who included 1D non-LTE corrections ( GLYPH<0> 0 : 01 dex for optical lines but GLYPH<0> 0 : 05 :: GLYPH<0> 0 : 10 for the IR lines) from Takeda et al. (2005) to their 3D LTE results. The value is unchanged from Asplund et al. (2009).  \nPotassium (Z = 19): The recommended solar K abundance in Asplund et al. (2009) was log GLYPH<15> K = 5 : 03 GLYPH<6> 0 : 09, which Scott et al. (2015b) updated to log GLYPH<15> K = 5 : 04 GLYPH<6> 0 : 05. The latter result was based on a 3D LTE analysis of five weak lines (excluding the strong 766.5 and 769.9 nm K i resonance lines) together with 1D non-LTE abundance corrections taken from Takeda et al. (1996) in the absence of consistent non-LTE calculations with a 3D solar model. Reggiani et al. (2019) recently presented a new 1D non-LTE analysis with a comprehensive model atom consisting of 133 atomic levels of K i plus the ground state of K ii and 250 radiative bound-bound transitions. Particular attention was placed on calculating or adopting reliable cross-sections for inelastic collisions, both with electrons, and with neutral hydrogen. Their 1D non-LTE calculations with a theoretical hydrostatic MARCS model atmosphere using the equivalent widths of three lines (including K i 769.9nm) yielded a solar abundance of log GLYPH<15> K = 5 : 11 GLYPH<6> 0 : 01 (standard deviation only).  \nHere, we present a full 3D non-LTE calculations for K i using the same model atom as in Reggiani et al. (2019). A sum-  \nmary of the inferred abundances for the di GLYPH<11> erent lines using a variety of model atmospheres is provided in Table A.1. Compared to Scott et al. (2015b), the gf -values have been updated to those recommended by NIST (Kramida 2019), which in all cases are estimated to have an accuracy better than 1%. Furthermore, the spectral line profiles were re-examined and the equivalent widths remeasured, leading us ultimately to exclude the 404.4 and 580.2 nm lines due to their substantial and uncertain blending. Our recommended solar K abundance is thus based on the measured equivalent widths of three K i lines: 693.9, 1177.0, and 1252.2 nm. Although we did not consider the K i 769.9 nm resonance line for the final result due to its large line strength and very pronounced non-LTE e GLYPH<11> ects ( GLYPH<25> + 0 : 2 dex), we note that its 3D non-LTE-based solar K abundance is fully consistent with those of the weak K i lines. Nevertheless, the departures from LTE are noticeable for the three retained K i lines (3D non-LTE - 3D LTE GLYPH<25> GLYPH<0> 0 : 05 dex), which reinforces the importance of a consistent 3D non-LTE analysis as employed here. Our recommended solar K abundance, based on the unweighted mean of the three K i lines, is log GLYPH<15> K = 5 : 07 GLYPH<6> 0 : 01 (statistical) GLYPH<6> 0 : 03 (systematic).  \nCalcium (Z = 20): Scott et al. (2015b) slightly revised the present-day photospheric Ca abundance of Asplund et al. (2009) to log GLYPH<15> Ca = 6 : 32 GLYPH<6> 0 : 03 with excellent agreement between the eleven Ca i and five Ca ii lines employed. These results were based on 3D LTE line formation calculations with our own 1D non-LTE abundance corrections calculated on the < 3D > model atmosphere, which range from + 0 : 02 to GLYPH<0> 0 : 03 dex for Ca i , and 0 : 00 to GLYPH<0> 0 : 07 dex for Ca ii . Recently, Mashonkina et al. (2017) and Osorio et al. (2019) presented new 1D non-LTE investigations that predict more severe, negative, non-LTE corrections for both species but more so for Ca ii , which would worsen the ionisation balance obtained in Scott et al. (2015b).  \nIn light of these issues, we carried out a detailed reanalysis of Ca, based on a consistent 3D non-LTE line formation study for the Sun. Our new Ca model atom consists of 141 Ca i levels, 46 Ca ii levels, and the Ca iii ground state, which are coupled radiatively by 2271 bound-bound and 135 bound-free transitions. We employed data for inelastic collisions with electrons from Zatsarinny et al. (2019b) for Ca i and from Meléndez et al. (2007) for Ca ii ; while for inelastic collisions with neutral hydrogen we draw on data from Barklem (2016b) and Belyaev et al. (2018) for Ca i and Ca ii , respectively. Osorio et al. (2020) reported a nonnegligible influence of non-LTE e GLYPH<11> ects on Mg i influencing the statistical equilibrium of Ca in the Sun. We therefore ran test calculations with the background Mg opacity treated in non-LTE; however, these were found to have a negligible impact on our results, at least for the Ca lines adopted in our abundance analysis.  \nAsummary of the inferred abundances for the di GLYPH<11> erent lines using a variety of model atmospheres is provided in Table A.1. Compared to Scott et al. (2015b), we adopted the same line data and mostly the same lines, dropping only the 616 : 13 nm that is a GLYPH<11> ected by the broad damping wings of the neighbouring Ca i 616 : 22 nm line. The unweighted mean results from the ten Ca i lines and the five Ca ii lines agree to better than 0 : 005 dex, with a very small line-to-line scatter for each species. Our recommended solar Ca abundance is log GLYPH<15> Ca = 6 : 30 GLYPH<6> 0 : 03 (total).", '2.5. Fe-peak elements: Scandium to nickel': "Scandium (Z = 21): Scott et al. (2015a) relied on a 3D LTE analysis of five Sc i and nine Sc ii lines coupled with non-LTE corrections computed with 1D model atmospheres to determine a solar Sc abundance of log GLYPH<15> Sc = 3 : 16 GLYPH<6> 0 : 04. With the new, highly  \naccurate experimental transition probabilities of Lawler et al. (2019), the weighted average becomes log GLYPH<15> Sc = 3 : 14 GLYPH<6> 0 : 01 (statistical) GLYPH<6> 0 : 04 (systematic) 6 . It is noted that the departures from LTE are substantial for the Sc i lines, but the weighted mean Sc abundance is largely driven by Sc ii . There is good agreement between the two ionisation stages, both of which give noticeably higher abundances than the CI meteoritic value (Sect. 3.1). Lawler et al. (2019) considered roughly twice as many Sc lines in their solar analysis, many of which are clean and have excellent atomic data; we postpone for later a study of these lines in conjunction with 3D non-LTE calculations and realistic collisional cross-sections.  \nTitanium (Z = 22): In their analysis of Fe-peak elements, Scott et al. (2015a) already benefited from the highly accurate experimental transition probabilities for Ti from the Wisconsin atomic physics group (Ti i : Lawler et al. 2013; Ti ii : Wood et al. 2013 7 ). They also applied non-LTE abundance corrections calculated with 1D model atmospheres, which are significant for Ti i ( GLYPH<25> + 0 : 06 dex) but unimportant for Ti ii . Scott et al. (2015a) found a large and unexplained di GLYPH<11> erence in Ti abundance from the two ionisation stages: 4 : 88 GLYPH<6> 0 : 05 versus 4 : 97 GLYPH<6> 0 : 04 from Ti i and Ti ii , respectively; their recommended solar Ti abundance was a weighted average of the two, resulting in log GLYPH<15> Ti = 4 : 93 GLYPH<6> 0 : 04. It is noteworthy that all of the considered theoretical and semi-empirical 1D model atmospheres implied significantly larger Ti i abundances than with the 3D Stagger model and more in line with the Ti ii results. This strong model atmosphere sensitivity suggests that the adopted 1D non-LTE corrections are not appropriate in 3D and that a consistent 3D non-LTE analysis would yield higher Ti i abundances. In contrast to Scott et al. (2015a), we based our recommended solar Ti abundance solely on Ti ii but increased the final error somewhat to reflect the discrepancy between the ionisation stages: log GLYPH<15> Ti = 4 : 97 GLYPH<6> 0 : 05 (total). A full 3D non-LTE study with realistic collisional cross-sections would be highly desirable; Lawler et al. (2013) and Wood et al. (2013) list a large number of highquality Ti lines with now excellent atomic data available.  \nVanadium (Z = 23): The recommended solar photospheric V abundance of Scott et al. (2015a) was based on a 3D LTE analysis of V i lines with an ad hoc non-LTE correction of 0.10 dex across the board in the absence of actual statistical equilibrium calculations based on the behaviour of Sc and Ti in 1D model atmospheres. As a minority species, V i is very sensitive to the atmospheric structure: the four considered theoretical and semiempirical 1D model atmospheres imply a 0 : 07 GLYPH<0> 0 : 18 dex higher abundance. Unfortunately, little guidance is available from the few available V ii lines due to their heavy blending, although nominally they suggest a GLYPH<25> 0 : 1 dex higher V abundance, both in 3D and 1D. Updating the Scott et al. (2015a) results with the Lawler et al. (2014) and Wood et al. (2014a) experimental gf -values, the mean V i abundance is only + 0 : 01 dex higher, while  \nthe V ii result changes by GLYPH<0> 0 : 02 dex 8 . We caution that our ad hoc V i non-LTE correction may well be underestimated, and we strongly encourage the first non-LTE investigation of this element to be performed. In the meantime we recommend the value found by Scott et al. (2015a) updated with the Lawler et al. (2014) transition probabilities: log GLYPH<15> V = 3 : 90 GLYPH<6> 0 : 08 (total).  \nChromium (Z = 24): The solar photospheric Cr abundance of Scott et al. (2015a) was based on the weighted average of a large number of Cr i and Cr ii lines with 1D non-LTE abundance corrections. Updating this result with the new experimental Cr ii transition probabilities of Lawler et al. (2017) leads to the same mean abundance: log GLYPH<15> Cr = 5 : 62 GLYPH<6> 0 : 04.  \nManganese (Z = 25): A new solar Mn abundance analysis was presented by Bergemann et al. (2019) based on full 3D non-LTE calculations. Using the same 3D hydrodynamical solar model (Magic et al. 2013) as for other elements and equipped with new ab initio inelastic Mn i + H i collisional cross-sections from Belyaev & Voronov (2017) (see also Grumer & Barklem 2020), they obtained log GLYPH<15> Mn = 5 : 52 GLYPH<6> 0 : 04. This is 0.10 dex larger than the result of Scott et al. (2015a)(5 : 42 GLYPH<6> 0 : 04) based on 3D LTE model spectra combined with < 3D > non-LTE abundance corrections; the two results do not overlap within the error bars. The two studies have very few Mn i lines in common, making a detailed comparison di GLYPH<14> cult. The abundance di GLYPH<11> erence cannot be traced to the adopted transition probabilities since the same sources are employed (mostly Den Hartog et al. 2011 and Blackwell-Whitehead & Bergemann 2007). Departures from LTE are obviously very prominent, with the 1D non-LTE GLYPH<0> 1D LTE di GLYPH<11> erence being + 0 : 07 dex. according to Bergemann et al. (2019). One may suspect that the 3D non-LTE e GLYPH<11> ects and Mn abundances in Bergemann et al. (2019) have been somewhat overestimated due to incomplete UV line blanketing over-estimating the photo-ionisation rates and thus weakening the predicted Mn i lines; a detailed comparison is further complicated by Bergemann et al. (2019) employing di GLYPH<11> erent non-LTE radiative transfer codes for their 1D and 3D analyses.  \nUntil the situation has been clarified, we advocate the recommended photospheric Mn abundance of Scott et al. (2015a) but with increased errors to reflect the remaining uncertainties associated with the non-LTE modelling: log GLYPH<15> Mn = 5 : 42 GLYPH<6> 0 : 06. Some confidence in this result can be had by noting that Scott et al. (2015a) obtained a value of 5 : 43 from their consistent < 3D > non-LTE analysis, only 0 : 01 dex larger. Our expectation is that, at least for atomic and ionic species in the disc-centre intensity spectrum of the Sun, the full 3D non-LTE result for Mn should be similar to the < 3D > non-LTE abundance, as we have found in our investigations of Li, C, N, O, Na, Mg, Al, Si, K, Ca, and Fe. Nevertheless, we encourage the issue to be revisited with new 3D non-LTE calculations with realistic inclusion of blending lines both for the statistical equilibrium and for the Mn i abundance diagnostic lines.  \nIron (Z = 26): In astronomy, Fe is often used as a proxy for the overall metallicity of a star on account of the element's relatively high abundance and the wealth of readily measurable lines. Iron is also an important opacity contributor in stellar interiors. In the Sun, many lines of both Fe i and Fe ii with accurate experimental atomic data are available for a spectroscopic analysis (e.g. Grevesse & Sauval 1999; Asplund et al. 2000c;  \nCa GLYPH<11> au et al. 2011; Mashonkina et al. 2011). From a 3D analysis, Scott et al. (2015a) obtained 7 : 45 GLYPH<6> 0 : 04 (including 1D non-LTE corrections of GLYPH<25> + 0 : 01 dex) and 7 : 51 GLYPH<6> 0 : 04, from Fe i and Fe ii lines, respectively, leading to a weighted average of log GLYPH<15> Fe = 7 : 47 GLYPH<6> 0 : 04 (statistical + systematic). Lind et al. (2017) updated this to log GLYPH<15> Fe = 7 : 47 GLYPH<6> 0 : 04 based on fully consistent 3D non-LTE calculations for Fe that for the first time were based on a large model atom that included ab initio collisional crosssections for Fe i + H i (Barklem 2018).  \nWe undertook a new analysis of Fe i and Fe ii using full 3D non-LTE calculations and improved atomic data. Our new Fe model atom consists of 100 Fe i levels, 76 Fe ii levels, and the Fe iii ground state, which are radiatively coupled through 4313 bound-bound and 100 bound-free transitions, amounting to 46000 wavelength points. We employed data for inelastic collisions with electrons from Wang et al. (2018) for Fe i and from Bautista et al. (2015) and Zhang & Pradhan (1995) for Fe ii ; while for inelastic collisions with neutral hydrogen, we drew on data from Barklem (2018) and Yakovleva et al. (2019) for Fe i and Fe ii , respectively. Photoionisation cross-sections for Fe i were sourced from Zatsarinny et al. (2019a). Recently, new accurate experimental transition probabilities for additional excellent solar Fe i lines have been published (Den Hartog et al. 2014; Belmonte et al. 2017), which we employed together with previous lines used by Scott et al. (2015a) with the transition probabilities recommended by the Oxford (e.g. Blackwell et al. 1995), Hannover (e.g. Holweger et al. 1991), and Wisconsin (e.g. O'Brian et al. 1991) groups 9 .  \nAsummary of the inferred abundances for the di GLYPH<11> erent lines using a variety of model atmospheres is provided in Table A.2. We now obtain excellent agreement between the two ionisation stages: 7 : 46 GLYPH<6> 0 : 04 and 7 : 47 GLYPH<6> 0 : 04 from Fe i and Fe ii respectively, based on unweighted means and including statistical and systematic uncertainties. Departures from LTE are small for both species in the Sun. There are no significant dependencies with wavelength, excitation potential, line strength or provenance of the gf -values in 3D, in contrast to the results using the 1D semiempirical Holweger & Müller (1974) model atmosphere (Fig. 2). With 1D theoretical models, the mean Fe abundance becomes unrealistically low (log GLYPH<15> Fe < 7 : 40), while with the < 3D > model, ionisation balance is not fulfilled, thus attesting to the need for full 3D non-LTE calculations. Our recommended photospheric Fe abundance is the average of the Fe i and Fe ii results: log GLYPH<15> Fe = 7 : 46 GLYPH<6> 0 : 04 (total), a value which has hardly moved over the past two decades in spite of major improvements in 3D and non-LTE analysis techniques and better atomic data (Asplund et al. 2000c, 2005a, 2009; Bergemann et al. 2012; Scott et al. 2015a; Amarsi et al. 2016b; Lind et al. 2017).  \nCobalt (Z = 27): We update the Scott et al. (2015a) results with the gf -values from Lawler et al. (2015) for the Co i lines in common: log GLYPH<15> Co = 4 : 94 GLYPH<6> 0 : 05 (total). While the new Co abundance is only 0 : 01 dex higher, the reduced line-to-line scatter is a testament to the high accuracy of the new experimental data. As for most Fe-peak elements, one concern relates to the adopted non-LTE corrections, which may be prone to error as they stem from our 1D calculations, and mainly draw on classical and semi-empirical recipes for inelastic collisions with electrons and with neutral hydrogen.  \n<!-- image -->  \nFig. 2. Inferred photospheric Fe abundances from detailed non-LTE calculations using the 3D hydrodynamical Stagger model ( Upper panel ) and the 1D Holweger & Müller (1974) model ( Lower panel ). The results for Fe i and Fe ii lines are shown as blue circles and red triangles, respectively. The solid lines denote the least-square fits to the Fe i results. In 3D non-LTE, there is excellent excitation and ionisation balance for Fe, contrary to the 1D case; it is noted that the trend with excitation potential with the Holweger & Müller (1974) model is even steeper in LTE (see Table A.2).  \n<!-- image -->  \nNickel (Z = 28): In the absence of any non-LTE calculations for Ni i , we adopt the 3D LTE based Ni abundance of Scott et al. (2015a), who used the accurate experimental transition probabilities from Wood et al. (2014b): log GLYPH<15> Ni = 6 : 20 GLYPH<6> 0 : 04. Based on our detailed 3D non-LTE calculations for Fe i , we expect the non-LTE e GLYPH<11> ects for Ni i to be similarly minor, but this requires confirmation.", '2.6. The heavy elements: Copper to uranium': "Relatively little work has been done recently to update the solar photospheric analysis by Grevesse et al. (2015) of the neutron capture elements. With the exception of the elements listed below, we adopt their recommended solar abundances. We note that many of these elements have few available spectral lines located in the crowded UV region, causing major problems in fitting the lines. Taking an even more conservative approach, we deemed some lines used by Grevesse et al. (2015) as too uncertain for a meaningful abundance determination even if very  \nfew, if any, alternative lines of the element are available. In some cases, we increased the error bars to attest to the challenges in working with heavily blended lines; often the blending lines are not even identified. The reader is furthermore cautioned that departures from LTE are largely unexplored territory for the heavy elements in the Sun. This deficiency can be expected to introduce significant systematic errors, quite possible at the level of GLYPH<6> 0 : 1 dex, in particular for many of the minority species.  \nCopper (Z = 29): We retain the solar Cu abundance log GLYPH<15> Cu = 4 : 18 GLYPH<6> 0 : 05 of Grevesse et al. (2015) despite it relying on old and relatively uncertain experimental oscillator strengths (e.g. Kock & Richter 1968) since the newer theoretical gf -values of Liu et al. (2014) imply an unacceptably large line-to-line abundance scatter. New laboratory data for this astrophysically important element would be welcome.  \nGermanium (Z = 32): We reinspected the spectral line profile of the quite strong but blended Ge i 326.95 nm line and revised the equivalent width downwards relative to Grevesse et al. (2015) to 4 : 20 pm. Our 3D LTE result consequently shifts downwards to log GLYPH<15> Ge = 3 : 62 GLYPH<6> 0 : 10.  \nRubidium (Z = 37): The only two available Rb i lines in the solar spectrum, the resonance lines at 780.03 and 794.76 nm, are unfortunately both perturbed by other species. We revised the measured line strengths of Grevesse et al. (2015) (to 0.53 and 0.29 pm, respectively) to better account for these blends and adjusted the remaining uncertainties to reflect the di GLYPH<14> culties in fitting the weak Rb i lines. The resulting 3D LTE abundance is slightly lower than before: log GLYPH<15> Rb = 2 : 44 GLYPH<6> 0 : 08 (total). Very recently, Korotin (2020) performed 1D non-LTE calculations for the Rb i resonance lines, employing ab initio data for inelastic collisions with neutral hydrogen (Yakovleva et al. 2018). They found substantial departures from LTE ( GLYPH<1> log GLYPH<15> Rb = GLYPH<0> 0 : 12 dex) and a solar Rb abundance consistent with the meteoritic value for the first time 10 . Their study employed blending lines in the statistical equilibrium computations as well as Rb spectrum synthesis. Applying the new non-LTE corrections of Korotin (2020) to our 3D LTE value above implies a solar Rb abundance of log GLYPH<15> Rb = 2 : 32 GLYPH<6> 0 : 08.  \nRhodium (Z = 45): We updated the Grevesse et al. (2015) Rh abundance with the new experimental transition probabilities for Rh i lines from Malcheva et al. (2015), while restricting the analysis to the 369.24 nm line due to the heavy blending of the 343.49 nm line. When also considering the uncertainties in the line strength and gf -values, we obtain log GLYPH<15> Rh = 0 : 78 GLYPH<6> 0 : 11 (statistical + systematic). The derived photospheric Rh abundance remains substantially below the value measured in CI chondrites, which may indicate pronounced departures from LTE for this minority species in the 3D model that are not yet possible to account for. It is noted that 1D semi-empirical solar atmosphere models yield abundances in better agreement with the meteoritic evidence; the shallower temperature gradient of the Holweger & Müller (1974) model tends to result in smaller departures from LTE (Asplund 2005), a phenomenon termed 'non-LTE masking' by Rutten & Kostik (1982).  \nPalladium (Z = 46): We reassessed the Pd i lines at 324.27 and 340.46 nm and revised their equivalent widths upwards rel-  \native to Grevesse et al. (2015), to 2 : 50 pm and 3 : 58 pm ; respectively. Our 3D LTE result consequently shifts upwards to log GLYPH<15> Pd = 1 : 57 GLYPH<6> 0 : 10, which is incidentally identical to the value recommended by Asplund et al. (2009).  \nBarium (Z = 56): Unfortunately, only a few spectral lines of Ba are available as abundance diagnostics, which furthermore are strong and saturated. The solar Ba abundances presented in Asplund et al. (2009) and Grevesse et al. (2015) were based on 3D LTE calculations to which 1D non-LTE abundance corrections were applied; naturally this approach is not fully selfconsistent, but it was the only one available at the time for this element. Very recently, Gallagher et al. (2020) revisited the issue using full 3D non-LTE computations with the same Stagger 3D hydrodynamical solar model as employed elsewhere in our study. Their inferred solar Ba abundance of log GLYPH<15> Ba = 2 : 27 GLYPH<6> 0 : 01 (statistical) GLYPH<6> 0 : 02 (systematic) was obtained from a 1D LTE spectrum synthesis with a MARCS model atmosphere (Gustafsson et al. 2008) to which abundance corrections from blending lines and 3D (i.e. the di GLYPH<11> erences in abundances in 3D non-LTE to obtain the 1D LTE equivalent width) were applied separately; this is not fully consistent either, but it is a more realistic approach than that in Grevesse et al. (2015). In addition to the three Ba ii lines employed by Grevesse et al. (2015), Gallagher et al. (2020) also considered the heavily blended 614.17 nm line and obtained consistent results. It is noted that the 3D LTE abundances in the two studies di GLYPH<11> er substantially for unknown reasons.  \nHere, we adopt the Gallagher et al. (2020) 3D non-LTE based result but inflate their uncertainties slightly to account for possible systematic errors from the model atmosphere, non-LTE calculations, and their assumed microturbulence in 1D: log GLYPH<15> Ba = 2 : 27 GLYPH<6> 0 : 05 (total). A new study directly comparing the 3D nonLTE line profiles (including blends) with observations would be welcome for this important element, especially given the di GLYPH<11> erence with the CI chondritic abundance (log GLYPH<15> Ba = 2 : 18 GLYPH<6> 0 : 02).  \nPraseodymium (Z = 59): We apply the 1D non-LTE abundance correction ( + 0 : 03 dex) for Pr i from Mashonkina et al. (2009) to the 3D LTE results of Grevesse et al. (2015), which in turn was based on 3D abundance corrections applied to the 1D spectrum synthesis of Sneden et al. (2009). The revised Pr abundance becomes log GLYPH<15> Pr = 0 : 75 GLYPH<6> 0 : 05.  \nTungsten (Z = 74): Compared to Grevesse et al. (2015), we remove the weak 484 : 38 nm line owing to uncertainties in measuring its equivalent width. Therefore, restricting the 3D LTE analysis of Grevesse et al. (2015) to the W i 400.875 nm line implies log GLYPH<15> W = 0 : 79 GLYPH<6> 0 : 11.  \nOsmium (Z = 76): After inspecting the spectral line profile of the Os i 330.16 nm transition, we revise its equivalent width downwards relative to Grevesse et al. (2015), to 0 : 86 pm. The resulting 3D LTE result is 0 : 05 dex lower: log GLYPH<15> Os = 1 : 35 GLYPH<6> 0 : 12.  \nIridium (Z = 77): We judge that the only detectable Ir i line, at 322.08 nm, is too badly blended and has too uncertain a continuum placement to allow us to make a meaningful solar Ir abundance measurement.  \nGold (Z = 79): Accounting for the uncertain continuum placement around the weak and badly blended Au i 312.28 nm line leads to a more conservative error estimate than in Grevesse et al. (2015): log GLYPH<15> Au = 0 : 91 GLYPH<6> 0 : 12 (total).  \nLead (Z = 82): We update the solar Pb abundance based on our new measurement of the equivalent width (0.91 pm) of the Pb i 368.35 nm, which results in log GLYPH<15> Pb = 1 : 95 GLYPH<6> 0 : 08, 0 : 03 dex larger than in Grevesse et al. (2015).  \n2.7. Elements observed in sunspots: Fluorine, chlorine, indium, and thallium  \nFluorine (Z = 9): While no spectral lines of F are detectable in the photospheric spectrum of the quiet Sun, there are features of the HF molecule in sunspot spectra that enable a determination of the solar F abundance. Maiorca et al. (2014) analysed the umbral atlas of Wallace et al. (2001) with a 1D theoretical model atmosphere corresponding to an e GLYPH<11> ective temperature of T e GLYPH<11> = 4250 K to estimate log GLYPH<15> F = 4 : 40 GLYPH<6> 0 : 25; most of the uncertainty stems from the choice of model atmosphere since the employed radiative-equilibrium model without magnetic fields may not be appropriate for this particular sunspot umbra. Here, we recommend the value of Maiorca et al. (2014), which represents a 0.16 dex reduction compared to the older Hall & Noyes (1969) result adopted by Asplund et al. (2009).  \nChlorine (Z = 17): Similar to F, the photospheric Cl abundance can only be estimated from sunspot spectra since the HCl molecule requires cool temperatures; no atomic features of Cl are present in the solar spectrum. Recently, Maas et al. (2016) revisited the previous HCl sunspot study of Hall & Noyes (1972) with improved molecular data and observations to obtain a 0.19 dex lower abundance: log GLYPH<15> Cl = 5 : 31 GLYPH<6> 0 : 12. Given the remaining model challenges, especially in finding a suitable model atmosphere for a sunspot of unknown T e GLYPH<11> , we argue that the quoted uncertainty is likely overly optimistic, and we instead adopt GLYPH<6> 0 : 20 dex. A more accurate determination of the solar abundances of the halogen elements would be possible through the abundance ratio of another species with similar temperature dependence and spectral features with well-determined gf -values, such as FeH and MgH.  \nIndium (Z = 49): Analyses of the In i 451.13 nm spectral line in the quiet Sun have traditionally yielded an In abundance much higher than the CI chondritic value (e.g. Bord & Cowley 2002: log GLYPH<15> In = 1 : 56 GLYPH<6> 0 : 20). Vitas et al. (2008) showed that this is likely due to an unknown blend and instead confirmed that the solar In abundance is close to the meteoritic value based on a sunspot spectrum: log GLYPH<15> In = 0 : 80 GLYPH<6> 0 : 20. We adopt this value here, as we did in Asplund et al. (2009).  \nThallium (Z = 81): The most recent determination of the solar Tl abundance dates back almost half a century. Updating with the current NIST-recommended transition probability for the only viable but heavily blended Tl i 535.04 nm line, the sunspot analysis of Lambert et al. (1972) suggests log GLYPH<15> Tl = 0 : 92 GLYPH<6> 0 : 17, quite likely with an underestimated uncertainty. A more accurate determination would be possible from the abundance ratio with another similar species such as Zr i .", '2.8. Indirect abundance determinations: Noble elements': 'The photospheric abundance of the noble gases cannot be inferred directly from spectroscopy due to the extremely high excitation potentials and thus minuscule level populations involved. Other indirect methods are therefore required.  \nHelium (Z = 2): The solar He abundance can be derived from the corona, prominences, solar wind, solar-energetic particles, and the gas giant planets. However, the He / Hratio varies greatly depending on environment, from 0.01 to 0.09 (e.g. Bochsler 2007; Reames 2014). The Genesis sample return mission of the solar wind implies a ratio of about 0.05 (Huss et al. 2020), which is significantly lower than expected in the Sun as a result of Coulomb-drag fractionation.  \nHere, we instead rely on helioseismology, which enables an accurate determination of the He abundance in the solar con-  \nvection zone and thus photosphere, which cannot be probed spectroscopically. The change in the adiabatic index GLYPH<0> 1 = ( @ ln P =@ ln GLYPH<26> ) S in the He ii ionisation zone occurring at temperatures T GLYPH<25> 10 5 K (corresponding to a radius of r GLYPH<25> 0 : 98 R GLYPH<12> ) leaves an observable signature in the p -mode oscillation frequencies that depends on the He content (e.g. Basu 2016 and references therein). In Asplund et al. (2009), we adopted the helioseismic He mass fraction value of Y surface = 0 : 2485 GLYPH<6> 0 : 0034 from Basu & Antia (2004); most of the quoted uncertainty stems from the choice of equation of state (OPAL: Rogers & Nayfonov 2002 vs. MHD: Mihalas et al. 1988). Due to the overlapping higher ionisation zones of other abundant elements, especially O, there is a degeneracy with the adopted metal mass fraction Z surface in the reference solar model for the helioseismic inversion: GLYPH<14> Y surface =GLYPH<14> Z surface = GLYPH<0> 1 : 06 and GLYPH<0> 0 : 56 with the OPAL (Basu & Antia 2004) and MHD (Trampedach et al. 2006) equation of states, respectively. Using a similar technique but with the improved SAHA-S3 equation of state, Vorontsov et al. (2013) obtained maximum-likelihood estimates of Y surface = 0 : 240 GLYPH<0> 0 : 255 and Z surface = 0 : 008 GLYPH<0> 0 : 013.  \nTaking the mean of the Basu & Antia (2004) and Vorontsov et al. (2013) He abundances corrected to our recommended present-day photospheric composition ( Z surface = 0 : 0139, Sect. 5) implies Y surface = 0 : 2423 GLYPH<6> 0 : 0054, which we adopt here. In terms of the standard logarithmic scale, this corresponds to log GLYPH<15> He = 10 : 914 GLYPH<6> 0 : 013, or He / H = 0.082 by number density.  \nNeon (Z = 10): The absolute Ne abundance can be determined from the radiation originating in the high-temperature regions of the Sun (e.g. transition region, corona, flares) or in situ measurements of the solar wind and solar energetic particles. Both types of analyses are, however, complicated by the existence of the still poorly understood first ionisation potential (FIP) e GLYPH<11> ect (e.g. Laming 2015); the abundances in the upper solar atmosphere and solar wind are modified relative to the photospheric values by a degree inversely dependent on the elemental ionisation potential. Arguably a more reliable approach, that we employ here, is to combine the abundance ratio of Ne ( GLYPH<31> ion = 21 : 56 eV) with an element with a similarly high ionisation potential and its photospheric abundance, under the assumption that the two elements experience the same, or at least a very similar, FIP e GLYPH<11> ect.  \nA common comparison element is O ( GLYPH<31> ion = 13 : 62 eV) since it can be measured in conjunction with Ne in the upper atmosphere and solar wind, and its photospheric abundance can be determined spectroscopically (albeit with challenges, Sect. 2.3); it is certainly conceivable that the di GLYPH<11> erence in ionisation potential could still introduce elemental variations. Young (2018) recently revisited the measured abundance ratios in the transition region of the quiet Sun, finding a significantly higher Ne / O ratio than previously estimated: Ne / O = 0 : 24 GLYPH<6> 0 : 05 compared with 0 : 175 GLYPH<6> 0 : 031 by Young (2005). Most of this adjustment comes from improvements in the atomic data, resulting in changes to the ionisation and recombination rates. Together with our photospheric O abundance, this would imply a photospheric Ne abundance of log GLYPH<15> Ne = 8 : 08 GLYPH<6> 0 : 09, a 40% increase compared with the value recommended in Asplund et al. (2009) using the Young (2005) Ne / O ratio.  \nLandi et al. (2007) inferred the absolute coronal Ne abundance from solar flare spectra: log GLYPH<15> Ne = 8 : 11 GLYPH<6> 0 : 12. There are, however, variations between flares both in the absolute Ne abundances (e.g. log GLYPH<15> Ne = 7 : 92 GLYPH<6> 0 : 03 Landi & Testa (2015) and in the Ne / O ratio during the solar cycle as a result of the FIP effect. It is expected that the results are the least susceptible to the FIP e GLYPH<11> ect during solar minima, for which Landi & Testa  \n(2015) obtained Ne / O = 0 : 25 GLYPH<6> 0 : 05, in perfect agreement with the quiet Sun transition region determination by Young (2018). The Ne / O ratio is typically lower in solar energetic particles (e.g. 0 : 140 GLYPH<0> 0 : 157 GLYPH<6> 0 : 010, Reames 2020) and solar wind (e.g. 0 : 11 GLYPH<0> 0 : 14, Reisenfeld et al. 2013), which suggests the existence of additional elemental fractionation. Thus, inferences of the absolute solar composition directly from solar wind measurements (e.g. von Steiger & Zurbuchen 2016) are likely prone to systematics and are probably unreliable (for further discussion see Serenelli et al. 2016).  \nAn independent probe of the Ne abundance is o GLYPH<11> ered by the measurements from the Genesis sample return mission (Burnett et al. 2019). Genesis o GLYPH<11> ers a comparison with He, which has a more similar ionisation potential ( GLYPH<31> ion = 24 : 59 eV) to Ne than O has. The quoted 4 He / 20 Ne and 20 Ne / 22 Ne ratios from the Genesis study of Heber et al. (2009) together with the helioseismic He / H ratio in the convection zone lead to a solar Ne abundance of log GLYPH<15> Ne = 8 : 14 GLYPH<6> 0 : 01. When accounting for correlations in the measured Ne / He and H / He ratios in di GLYPH<11> erent solar wind regimes, the abundance becomes log GLYPH<15> Ne = 8 : 06 GLYPH<6> 0 : 03 (Burnett et al. 2019; Huss et al. 2020), which is in excellent agreement with the above estimate using the Young (2018) Ne / O ratio. In the absence of a proper understanding of the physical processes responsible for the FIP e GLYPH<11> ect, it is di GLYPH<14> cult to estimate the full uncertainties attached to this Genesis abundance. We tentatively assign half the di GLYPH<11> erence between Heber et al. (2009) and Burnett et al. (2019) as a systematic error, which leads to our final recommended solar photospheric Ne abundance of log GLYPH<15> Ne = 8 : 06 GLYPH<6> 0 : 03 (statistical) GLYPH<6> 0 : 04 (systematic). The excellent agreement with the coronal value by Young (2018) lends strong support to this 0.13 dex higher abundance than in Asplund et al. (2009) and consequently also the here advocated photospheric O abundance.  \nArgon (Z = 18): Asplund et al. (2009) used the straight average from a variety of methods to arrive at a solar Ar abundance of log GLYPH<15> Ar = 6 : 40 GLYPH<6> 0 : 13 following the procedure outlined by Lodders (2008): solar wind, solar flares and energetic particles, nuclear statistical equilibrium, Jupiter, and solar neighbourhood. Updating this exercise with new data and down-weighting the planetary and stellar data would lead to log GLYPH<15> Ar = 6 : 36 GLYPH<6> 0 : 11.  \nHere, we instead opt for a single, high-precision measurement: the Genesis sample return of the solar wind (Burnett et al. 2019), in the same manner as our recommended Ne abundance was determined. Allowing for correlations between the measured H / He and 36 Ar / 4 He ratios, Burnett et al. (2019) and Huss et al. (2020) obtained log GLYPH<15> Ar = 6 : 38 GLYPH<6> 0 : 08 (statistical). We estimate that the remaining systematic errors are significant given the still poorly understood FIP e GLYPH<11> ect (Laming 2015) and the differences in ionisation potential between Ar ( GLYPH<31> ion = 15 : 76 eV) and the reference element He ( GLYPH<31> ion = 24 : 59 eV); it is noted that Meshik et al. (2020) have argued that the FIP fractionation still varies with ionisation potential even for GLYPH<31> ion > 13 eV. We adopt log GLYPH<15> Ar = 6 : 38 GLYPH<6> 0 : 10 for the Sun.  \nKrypton (Z = 36): Here, we base the recommended solar Kr abundance on the Genesis measurements of the solar wind by Meshik et al. (2020). With the quoted Kr and H fluences, one obtains log GLYPH<15> Kr = 3 : 12 GLYPH<6> 0 : 03 (statistical). This value is preconditioned on the assumption that the FIP fractionation corrections are identical for Kr ( GLYPH<31> ion = 14 : 00 eV) and H ( GLYPH<31> ion = 13 : 60 eV) even when individually the corrections are about a factor of three. Alternatively, using the Kr and Ar ( GLYPH<31> ion = 15 : 76 eV) fluences together with the above estimated solar Ar abundance leads to a slightly higher value with a greater statistical uncertainty: log GLYPH<15> Kr = 3 : 16 GLYPH<6> 0 : 08. Burnett et al. (2019) found a Kr  \nabundance 0.2 dex lower with the same Genesis data when using O ( GLYPH<31> ion = 13 : 62 eV) as the reference element, which may imply element-specific corrections for high-FIP elements beyond theoretical predictions (Meshik et al. 2020), or that the true photospheric O abundance is correspondingly much higher than advocated here (Laming et al. 2017). In all likelihood, the uncertainty is dominated by systematic errors from the FIP e GLYPH<11> ect, whichever method is chosen. We adopt log GLYPH<15> Kr = 3 : 12 GLYPH<6> 0 : 10 including estimated systematic errors.  \nTraditionally, the solar Kr abundance has been estimated from the smooth variations in the abundances of nearby oddmass number isotopes in CI chondrites (Sect. 3) as predicted by the measured cross-sections for the slow neutron capture process (e.g. Wisshak et al. 1998), a method originally proposed by Suess (1947a,b). With updated isotopic abundances, this leads to log GLYPH<15> Kr = 3 : 15 GLYPH<6> 0 : 10 (statistical). The true uncertainty is likely larger given that the odd-nuclei abundance smoothness has been questioned and that a more realistic distribution of neutron exposures than in this classical model leads to significantly di GLYPH<11> erent predictions for Kr (Arlandini et al. 1999). Nevertheless, it is reassuring that this completely independent method agrees well with the estimate based on the solar-wind.  \nXenon (Z = 54): The Xe abundance can be inferred via the photospheric Sm measurement combined with experimental neutron capture cross-sections under the assumption that the reaction flow has reached equilibrium. Using their measured cross-sections for the pure s-process isotopes, 130 Xe and 150 Sm Reifarth et al. (2002) obtained log GLYPH<15> Xe = 2 : 24 GLYPH<6> 0 : 01 (statistical). With our updated photospheric Sm abundance, we revise this estimate slightly to log GLYPH<15> Xe = 2 : 22 GLYPH<6> 0 : 05, including estimated systematic errors (adopting the CI chondritic Sm abundance instead to normalise the s -process production would imply a 0.01 dex lower Xe abundance). The solar wind abundance inferred from Genesis data (Meshik et al. 2020) is significantly higher, log GLYPH<15> Xe = 2 : 42 GLYPH<6> 0 : 05 (statistical), suggesting a di GLYPH<11> erent FIP fractionation between Xe ( GLYPH<31> ion = 12 : 13 eV) and H, as also expected on theoretical grounds (Laming et al. 2017).', '3.1. Chemical composition of CI chondrites': 'Independent and complementary data on the overall chemical composition of the Solar System are o GLYPH<11> ered by meteorites. The vast majority of the more than 50,000 meteorites found on Earth have undergone fractionation and alteration to varying degrees, leaving just a handful of the most primitive and undi GLYPH<11> erentiated carbonaceous chondrites, the so-called CI chondrites: the Alais, Ivuna, Orgueil, Revelstoke, and Tonk meteorites, named after the locations of the falls 11 . On account of the overall similarities with the solar photospheric abundances, it has long been argued that the CI chondrites contain pristine samples of the original proto-solar nebula (e.g. Urey 1952, Cameron 1968; see Lodders 2019 for a historical overview). The aqueous alteration and thermal metamorphism of the textural and mineralogical structure all CI chondrites have experienced are not believed to have significantly changed its overall chemical composition, with the exception of the highly volatile elements that are strongly depleted  \nFig. 3. Comparison between the recommended CI chondritic abundances by Palme et al. (2014), which we have adopted here, with other recent CI analyses (red symbols: Barrat et al. 2012; blue: Braukmüller et al. 2018; turquoise: King et al. 2020; green: Lodders (2019). Circle symbols correspond to data for the Orgueil meteorite, star symbols are for the Ivuna CI chondrite, and triangles for mean of all CI chondrites. All data are published abundances (ppm by mass) relative to Palme et al. (2014).  \n<!-- image -->  \n(Palme et al. 2014) 12 . Since these depleted elements are also the most cosmically abundant, the meteoritic abundances by number densities are therefore traditionally normalised relative to Si: N Si GLYPH<17> 10 6 (Suess & Urey 1956). To anchor the two sets of abundances onto the same scale, the meteoritic and photospheric Si abundances are therefore set equal with our recommended solar Si abundance (Sect. 2.4): log GLYPH<15> X = 1 : 51 + log N X. In principle, more elements can be included in the conversion (e.g. Anders & Grevesse 1989; Lodders et al. 2009; Palme et al. 2014, who used 12, 39, and 38 elements, respectively), but for simplicity we only employ Si, which was also done by Asplund et al. (2005a, 2009) and Lodders (2019); we would have obtained the same result had we adopted the same normalisation elements as Palme et al. (2014).  \nSeveral literature compilations of the chemical compositions of CI chondrites exist, with Wasson & Kallemeyn (1988), Anders & Grevesse (1989), Lodders (2003), Lodders et al. (2009), Palme et al. (2014), and Lodders (2019) being particularly highly regarded and authoritative. These studies provide a comprehensive curation of a large number of original references of laboratory data on meteorites. Here (see Table 2), we adopt the CI chondritic abundances recommended by Palme et al. (2014) for the 83 elements for which data from the Orgueil meteorite exist. To convert the meteoritic abundance measurements by mass to the standard meteoritic abundance scale by number, we used the atomic weights recommended by Meija et al. (2016a) with a few adjustments to accommodate more representative Solar System isotopic ratios as described in Sect. 4. The quoted CI abundance uncertainties are only statistical sampling errors.  \nWe did not attempt to factor in the recent measurements by Braukmüller et al. (2018) on account of the surprisingly large variation between their analyses of Orgueil and Ivuna, as well as  \nFig. 4. Comparison between our recommended abundances in presentday solar photosphere and CI chondrites. Most elements agree well, with the exception of the highly volatile elements, which are depleted in meteorites, and Li, which has been largely destroyed in the Sun due to nuclear burning.  \n<!-- image -->  \nwith prior CI measurements for several elements, including Na, K, Zr, the rare Earth elements, and Th and U (Fig. 3). The Ivuna analysis by King et al. (2020) agrees well with the Palme et al. (2014) recommended abundances overall, although some striking discrepancies exist, including for S, P, K, Zr, Ba, Hf, Ta, W, and Th and U (Fig. 3). These di GLYPH<11> erences can largely be traced to increased mobility in fluids for some elements during the aqueous alterations of the CI chondrites (Palme et al. 2014; Braukmüller et al. 2018). The recommended values of Palme et al. (2014) agree very well with the newer compilation by Lodders (2019), with a few exceptions. Notable changes have occurred for C, N, F, Br, Sb, I, and Hg, as seen in Fig. 3, but those elements are also among the most uncertain. Lodders (2019) performed a weighted average of the available analyses of all CI chondrites, but since they did not provide references to the original data, we instead opted for the compilation of Palme et al. (2014).', '3.2. CI chondritic and photospheric abundances': 'Figure 4 shows a comparison of our recommended present-day solar photospheric abundances with those from CI chondrites on the astronomical logarithmic abundance scale. Clearly, the agreement is very good overall, with a few striking and wellknown di GLYPH<11> erences. The noble gases (He, Ne, Ar, Kr, Xe) are highly depleted in CI chondrites by factors ranging from > 10 4 (Xe) to > 10 9 (He, Ne), while the other highly volatile gases H, C, N, and O are depleted by factors of 10 4 , 10, 40, and 2, respectively.  \nThe solar photospheric Li abundance is a factor of almost 200 lower than in CI chondrites due to nuclear processing in the Sun, much more than predicted by standard solar models. It implies additional mixing below the solar convection zone, bringing gas down to high enough temperatures ( GLYPH<25> 2 : 5 GLYPH<1> 10 6 K) for  \nnuclear destruction to take place and subsequently returning it to the convection zone (e.g. Charbonnel & Talon 2005). Several possible physical processes for this mixing have been proposed (e.g. internal gravity waves, rotation, convective overshooting, and atomic di GLYPH<11> usion or a combination thereof), but all available modelling still requires parametrisation and hence remains uncertain. Solar twins show a slow and monotonically increasing Li depletion with age (Carlos et al. 2019), with the Sun being unusually deficient in Li by GLYPH<25> 0 : 3 dex given its age for reasons not yet understood. There is a long and chequered history surrounding the question of whether Li depletion in the Sun and similar stars is related to planet formation (e.g. Bouvier 2008; Israelian et al. 2009; Baumann et al. 2010; Bara GLYPH<11> e & Chabrier 2010).  \nGiven the existence of additional mixing below the convection zone as evident from Li, the photospheric abundances of Be and B take on particular importance, since they burn at slightly higher temperatures than Li ( GLYPH<25> 3 : 5 GLYPH<1> 10 6 K and GLYPH<25> 5 GLYPH<1> 10 6 K, respectively). The fact that solar twins spanning a wide age range all have very similar photospheric Be abundances suggests that the Sun has not experienced significant Be depletion in the convection zone, even though the solar Be abundance is perhaps GLYPH<25> 0 : 05 dex lower than expected for its age (Tucci Maia et al. 2015). A modest B deficiency could, in principle, be hidden within the large error bars, but the absence of a significant Be depletion argues against it. The photospheric Li, Be, and B abundances imply shallow extra mixing below the convection zone present over Gyr timescales in the Sun.  \nA more detailed comparison of the photospheric and CI chondritic abundances as a function of atomic number is shown in Fig. 5. With the exception of the depleted elements mentioned above, essentially all elements agree within their combined uncertainties. Using the inverse of the combined photospheric and meteoritic abundance uncertainties in quadrature as weights, the mean average di GLYPH<11> erence of the 58 non-depleted elements (i.e. excluding Li and the highly volatiles) amounts to 0 : 000 GLYPH<6> 0 : 007 (s.d. = 0 : 052) dex. We remind the reader that the meteoritic abundance scale has been normalised to the photospheric Si abundance, which is here only quoted to within 0.01 dex, and hence this perfect agreement is somewhat fortuitous. This comparison does, however, support our choice of transition probabilities for Si (Sect. 2.4). Likewise, it is noted that atomic di GLYPH<11> usion and gravitational settling taking place in the Sun over the past 4.57 Gyr will retain this overall agreement since all elements bar H and He di GLYPH<11> use in nearly identical amounts (Sect. 5). Of these 58 elements, four are nominal > 2 GLYPH<27> outliers: Sc, Rh, Ag, and Hf; we suspect that the photospheric abundances from the minority ionisation stages Rh i and Ag i to have been significantly underestimated in the absence of non-LTE calculations (Sect. 2.6). Overall, CI chondrites closely resemble the solar photospheric abundances, justifying their historical use as a compositional proxy. Caveats are in place, however, as discussed below.', '3.3. Chemical signatures of planet formation': 'The overall impressive consistency between the CI chondrites and the solar photosphere notwithstanding, it is worthwhile to ask whether the two abundance sets are truly the same. Or in other words, are CI chondrites a true reflection of the pristine chemical composition of the proto-solar nebula as has long been assumed (e.g. Suess & Urey 1956)? Figure 6 shows the differences between the photospheric and the CI meteoritic logarithmic abundances as a function of 50% equilibrium condensation temperature ( T cond) computed for solar composition gas  \nFig. 5. Di GLYPH<11> erences of the photospheric and CI chondritic logarithmic abundances as functions of atomic number. In most cases, the elemental abundances agree within the combined uncertainties; the highly volatile elements and Li are not shown in the figure due to meteoritic and solar depletion, respectively. The two most discrepant elements are Rh and Ag, both of which are suspected to have underestimated photospheric abundances in the absence of non-LTE calculations.  \n<!-- image -->  \nat a pressure of 10 GLYPH<0> 4 bar (Wood et al. 2019, see also Lodders 2003); for clarity, only elements with a combined abundance uncertainty of < 0 : 10 dex are included. With the anchoring of the two abundance scales through Si ( T cond = 1314 K), it appears as though the moderately volatile elements (300 < GLYPH<24> T cond < GLYPH<24> 1250 K; here: Pb, S, Zn, Rb, K, Ga, Cu, Na, Mn, in order of increasing T cond) are enriched in the CI chondrites relative to the Sun: the weighted mean di GLYPH<11> erence is GLYPH<0> 0 : 042 GLYPH<6> 0 : 006 (s.d. = 0 : 018) dex. Not surprisingly, the main component elements (1250 < GLYPH<24> T cond < GLYPH<24> 1370 K; here: Mg, P, Si, Cr, V, Fe, Co, Ni) agree excellently, due to the fact that we enforced the Si abundances to agree: + 0 : 004 GLYPH<6> 0 : 009 (s.d. = 0 : 025) dex. The refractory elements ( T cond > GLYPH<24> 1370 K, here 24 elements) may be slightly depleted in CI chondrites, but the relatively large scatter prevents a firm conclusion: + 0 : 015 GLYPH<6> 0 : 010 (s.d. = 0 : 048) dex. We caution that the absolute level of enhancement / depletion depends on the chosen reference element. It is well known that the CI chondrites are volatile rich compared to the other carbonaceous chondrites, but here we thus find evidence that this is also the case in comparison to the Sun. The opposite is of course true for the highly volatile ( T cond < GLYPH<24> 250 K) elements (H, He, C, N, O, Ne, Ar, Xe, Kr), which are greatly depleted in all meteorites, including the CI chondrites (Sect. 3.1). This disposition for volatiles and refractories in the Sun and CI chondrites was first noticed by Gonzalez (1997) when discussing the metalrichness of the first discovered exoplanet host stars, and it is here substantiated by our more accurate photospheric abundances.  \nThere is significant scatter in Fig. 6 around any possible trend, mostly driven by the photospheric abundance uncertainties; higher precision is naturally obtained when instead intercomparing relative abundances of similar objects analysed homogeneously, such as solar / stellar twins or chondrites. A weighted linear least-squares fit implies GLYPH<1> log GLYPH<15> = (6 : 9 GLYPH<6> 2 : 7) GLYPH<1> 10 GLYPH<0> 5 T cond + ( GLYPH<0> 0 : 09 GLYPH<6> 0 : 04) ; and hence the correlation does not quite reach 3 GLYPH<27> significance; with the condensation temperatures of Lodders (2003), the slope would be (7 : 3 GLYPH<6> 2 : 6) GLYPH<1> 10 GLYPH<0> 5 dex K GLYPH<0> 1 . Including all elements except Li and the highly volatile elements  \nhardly changes the best fit. The quality of the current abundance data does not motivate a more sophisticated function, such as piece-wise linear and step functions, as attempted in related cases involving dust condensation (e.g. Meléndez et al. 2009; Gonzalez 2014).  \nCompared to CI chondrites, other types of carbonaceous chondrites show increasing fractionation with condensation temperature (CI ! CM ! CR ! CO ! CV ! CB ! CH, Wasson & Kallemeyn 1988; Krot et al. 2014). Since Palme et al. (2014) did not provide data for other classes of chondrites, we made use of the published abundances by Bland et al. (2005) for 42 elements, re-normalised and converted to the astronomical logarithmic abundance scale. An advantage with the Bland et al. (2005) dataset is that they tabulate the compositions both for the bulk of the chondritic material as well as the fine-grained matrix of the meteorites, which are important for our discussion below. The abundances for other chondrite classes are given normalised to CI (and the element Yb). While the relative Si abundances are provided for the matrix components, they are missing for the bulk material for non-CI chondrites for reasons unclear to us. As a result, when comparing the bulk compositions we have had to normalise to the astronomical abundance scale using Mg as the reference element, which introduces a 0.02 dex shift; as will become clear, this minor o GLYPH<11> set is unimportant for our conclusions.  \nThe resulting CI abundances from Bland et al. (2005) are not identical to those adopted in Sect. 3.1, but the di GLYPH<11> erences are in most cases minor, as seen in Fig. 7. Importantly, the correlation with condensation temperature is still apparent even if data for some elements, including the moderately volatile Na and S, are lacking; the weighted di GLYPH<11> erence between the Sun and the Bland et al. (2005) CI data for the moderately volatile elements is GLYPH<0> 0 : 040 GLYPH<6> 0 : 007 (s.d. = 0 : 018) dex, which is almost identical to the case with the Palme et al. (2014) CI abundances adopted elsewhere in our study. In terms of their bulk abundances, the CI chondrites are clearly the most solar-like, the trend with condensation temperature notwithstanding, followed by the CM and CO meteorites, which show increasing depletion of the volatile elements, while the refractory elements display less of a variation relative to the Sun and CI (Fig. 7); the highly volatile elements (H, C, N, O, and noble gases) not included in the figure are even more depleted in the non-CI chondrites.  \nAn important distinction between the chondrite classes is that CI meteorites are only comprised of the fine-grained matrix, while other chondrites also consist of chondrules and refractory inclusions with a volatile-depleted / refractory-enriched composition relative to their matrices; the fraction of matrix material varies between the classes as well as between meteorites of the same type. Chondrules have experienced rapid heating to temperatures > 1400 K, followed by similarly quick cooling, likely in the inner part of the proto-planetary disc, but the details of their formation and subsequent incorporation into the matrix is still hotly debated; the matrix material has not experienced any significant heating events. The chemical complementarity between the chondrules and matrix that has been previously argued for (e.g. Bland et al. 2005; Palme & Hezel 2011) has more recently been questioned, implying that the chondrules and matrix did not necessarily form from the same reservoir (van Kooten et al. 2019; Alexander 2019). To facilitate a fair comparison with the solar abundances, one should therefore consider the chondritic matrices, which are clearly more pristine than the chondrules, rather than their bulk compositions.  \nThe lower panel of Fig. 7 compares the matrix compositions of CI, CM, and CO chondrites relative to the Sun using data from Bland et al. (2005). Since the chondritic matrix con-  \nFig. 6. Di GLYPH<11> erences of the photospheric and CI chondritic logarithmic abundances as a function of 50% condensation temperature at a pressure of 10 GLYPH<0> 4 bar for a solar composition gas (Wood et al. 2019), showing a tendency for the refractory elements to be enhanced in the Sun compared with the meteorites. Only elements with a combined uncertainty of < 0 : 1 dex in the abundance di GLYPH<11> erence are shown in the figure. It should be noted that C, N, and O (condensation temperature < 200 K) fall outside the plot with the meteoritic abundances depleted compared to the Sun, as is also the case for the noble gases. The solid line is a weighted least-squares fit to the data with a slope of (6 : 9 GLYPH<6> 2 : 7) GLYPH<1> 10 GLYPH<0> 5 dex K GLYPH<0> 1 .  \n<!-- image -->  \ntains a higher abundance of volatile elements than chondrules, the matrices are more similar to the Sun. Remarkably, the CM matrix is overall a very good match to the present-day solar photosphere, closer in fact than the CI chondrites. In particular, the moderately volatile elements agree better with a weighted mean di GLYPH<11> erence of + 0 : 014 GLYPH<6> 0 : 010 (s.d. = 0 : 027) dex, while the refractory elements show similar di GLYPH<11> erences for CI and CM matrix although of opposite sign; in both cases, the main component elements agree excellently with the Sun. The weighted linear least-squares fit for the CM matrix data has a slope of ( GLYPH<0> 4 : 6 GLYPH<6> 3 : 6) GLYPH<1> 10 GLYPH<0> 5 dex K GLYPH<0> 1 , while the corresponding slope for CI chondrites is (8 : 6 GLYPH<6> 3 : 6) GLYPH<1> 10 GLYPH<0> 5 dex K GLYPH<0> 1 with the Bland et al. (2005) elemental abundances. We conclude that CI chondrites are not a perfect compositional match with the Sun, as has been assumed for more than half a century. Instead, the solar abundances more closely resemble those of the matrices of CM chondrites, although not the bulk CM composition. This could reflect spatially separate origins of the CM matrix and chondrules with di GLYPH<11> ering chemical make-up, with some chemical complementarity possibly emerging later following transport to where amalgamation occur (van Kooten et al. 2019). A more detailed comparison will be required to elucidate how, where and when the di GLYPH<11> erent chondrite constituents formed and their relation to the solar composition in this new scenario (e.g. Scott 2007; Alexander 2019; van Kooten et al. 2019). In the meantime, we advise against using meteoritic abundances to estimate the Solar System composition, which is often done due to their generally higher precision (e.g. Vinyoles et al. 2017; Lodders 2019; Wang et al. 2019).  \nThe correlations with chondrites shown in Fig. 7 are reminiscent of the peculiar solar abundances relative to solar twins; the Sun is unusually volatile rich and refractory poor with an abundance amplitude of about 0 : 08 dex (Meléndez et al. 2009; Ramírez et al. 2010). In other words, the abundance di GLYPH<11> erences relative to the average of solar twins is almost exactly a mirror  \n<!-- image -->  \n1800  \n1800  \nFig. 7. Upper panel: Di GLYPH<11> erence between the solar photospheric and bulk abundances for CI (blue circles), CM (red triangles), and CO (green stars) chondrites as measured by Bland et al. (2005). CM and CO chondrites show increasing fractionation compared to the Sun, especially for the volatile elements. Lower panel: Same, but for the finegrained matrix of chondrites. Surprisingly, the solar abundances more closely resemble the CM matrix than the CI equivalent, which shows a trend with condensation temperature. As explained in the text, for the bulk comparison Mg has been used as the reference element to connect the meteoritic and photospheric abundance scales since the Si abundances are not provided by Bland et al. (2005) for non-CI chondrite classes, which introduces a 0.02 dex o GLYPH<11> set. For the matrix compositions, the usual reference element Si has been used.  \n<!-- image -->  \nimage of the di GLYPH<11> erences with the CI chondrites. Only GLYPH<25> 10% of the studied stars share the detailed chemical fingerprint with the Sun, which furthermore cannot be explained by di GLYPH<11> erences in age or Galactic birthplace of the stars (Bedell et al. 2018; Nissen &Gustafsson 2018). The correlation with condensation temperature suggests that the signature is related to planet formation. The initial suggestion that the abundance pattern is imprinted by the formation of terrestrial planets in the Solar System (Meléndez et al. 2009; Chambers 2010) is likely not correct since the solar convection zone was much deeper during the proto-planetary disc phase compared to today, thus requiring much more than a few M GLYPH<8> of refractory-rich material locked up in rocky planets. Furthermore, transit data from Kepler imply that most solarlike stars harbour super-Earths or mini-Neptunes more massive than the terrestrial planets in the Solar System (Zink et al. 2019),  \nwhich would make the Sun stand out in the opposite manner compositionally. The alternative explanation that other stars have on average experienced more accretion of rocky planets is also unlikely to hold up due to the rapid erasure of such metal-rich accretion from the stellar surface convection zone due to thermohaline mixing (Théado & Vauclair 2012).  \nRecently, a more promising scenario was proposed by Booth & Owen (2020), which involves the early formation of giant planets trapping > 100 M GLYPH<8> of refractory-rich dust external to their orbits while volatile-rich gas continues to accrete onto the proto-star. They found that such gas-dust separation could imprint an abundance pattern similar to that observed when the Sun is deficient in refractory elements by GLYPH<25> 0 : 08 dex (20%). The prediction is thus that most solar twins lack a gas giant planet, allowing the gas and dust accretion to continue largely unimpeded.  \nThe incidence of the peculiar solar abundance pattern is consistent with the frequency of giant planets being GLYPH<25> 7% (Wittenmyer et al. 2020). Long-term precision radial velocity monitoring of both stars most similar and most chemically disparate to the Sun should be undertaken to search for Jupiter-like planets; HIP11915 is also a chemically solar twin that hosts a Jupiter twin (Yana Galarza et al. 2021), although the planet hosts HIP5301 and HIP15227 seemingly do not fit the pattern (Bedell et al. 2018) - a much larger sample is obviously needed. The connection with the chondritic compositions discussed above also needs to be deciphered within the framework of an early formation of Jupiter.', '4. Isotopic abundances': "Isotopic abundances can only be measured from solar spectroscopy for C and O, and even then only with significant uncertainties (Scott et al. 2006; Ayres et al. 2013; Lyons et al. 2018) 13 . Instead, alternative methods are used to infer the solar values, including sample return missions of the solar wind such as NASA's Genesis probe (Burnett et al. 2003; Jurewicz et al. 2003) and laboratory measurements of various terrestrial, lunar, meteoritic, cometary, and asteroid samples. Our advocated isotopic abundances can be found in Table B.1.  \nThe present-day solar D abundance can not be directly determined due to near-complete burning of D to 3 He during the pre-main-sequence evolution. The increase from the proto-solar 3 He / 4 He ratio as indicated by the value measured in Jupiter (1 : 66 GLYPH<6> 0 : 05 GLYPH<1> 10 GLYPH<0> 4 , Maha GLYPH<11> y et al. 1998) and the local interstellar cloud (1 : 62 GLYPH<6> 0 : 29 GLYPH<1> 10 GLYPH<0> 4 , Busemann et al. 2006), to the presentday ratio in the solar wind is thus a reflection of the proto-solar D / H ratio. The isotope fluences determined from the Genesis probe suggest 3 He / 4 He = (4 : 5 GLYPH<0> 4 : 8) GLYPH<1> 10 GLYPH<0> 4 depending on the solar wind regime (Heber et al. 2012), which indicates the presence of isotopic fractionation during the formation and acceleration of the solar wind. Two methods to account for fractionation are presented in Heber et al. (2012): theoretical corrections based on the ine GLYPH<14> cient Coulomb drag model or those based on empirical correlations between 3 He / 4 He and H / 4 He in the solar wind, which is similar to our determination of the solar Ne and Ar absolute abundances. Substantially di GLYPH<11> erent present-day 3 He / 4 He ratios and thus proto-solar D / H abundances are obtained: 1 : 24 GLYPH<1> 10 GLYPH<0> 5 (theoretical) and 1 : 67 GLYPH<1> 10 GLYPH<0> 5 (empirical). We take the average of the two values together with the Jovian D / H(2 : 1 GLYPH<1> 10 GLYPH<0> 5 , Lellouch et al. 2001) to obtain a proto-solar D / Hratio of (1 : 67 GLYPH<6> 0 : 25) GLYPH<1> 10 GLYPH<0> 5 .  \nMass number  \nFig. 8. Proto-solar isotopic abundances as a function of mass number after taking into account atomic di GLYPH<11> usion for He ( + 0 : 07 dex) and heavier elements ( + 0 : 06 dex) and radioactive decay. Even isotopes are shown as filled blue circles and odd isotopes as open red circles.  \n<!-- image -->  \nWecaution that the here-derived proto-solar D / Hvalue is in contradiction with the present-day D abundance in the local interstellar medium of D / H GLYPH<21> (2 : 0 GLYPH<6> 0 : 1) GLYPH<1> 10 GLYPH<0> 5 (Linsky et al. 2006; Prodanovi'c et al. 2010) since astration of D in stars should result in a monotonic decrease with time. The here-deduced proto-solar D would also lead to uncomfortably large astration factors from the primordial value of (2 : 53 GLYPH<6> 0 : 03) GLYPH<1> 10 GLYPH<0> 5 as inferred from highredshift, metal-poor damped Ly-alpha systems (Weinberg 2017; Cooke et al. 2018). The proto-solar D / H isotopic ratio clearly deserves further scrutiny.  \nThe 12 C / 13 C ratio does not appear to vary significantly between di GLYPH<11> erent Solar System objects. We therefore adopt the reference terrestrial value of 89 : 3 GLYPH<6> 0 : 2 (Meija et al. 2016b), which is consistent with the solar photospheric determination (Scott et al. 2006; Lyons et al. 2018) and the preliminary Genesis results for the solar wind (Wiens et al. 2013).  \nThe 14 N / 15 N ratio varies tremendously between Solar System sources for reasons that are still poorly understood. The Genesis team has measured the isotopic ratio in the solar wind (Marty et al. 2010), which with an isotopic fractionation correction (Laming et al. 2017) implies 14 N / 15 N = 453 GLYPH<6> 66. A similar value, 435 GLYPH<6> 57, is obtained for Jupiter (Owen et al. 2001). We take the weighted mean of these two measurements: 14 N / 15 N = 443 GLYPH<6> 13.  \nThe isotopic composition of O in the solar wind was measured by McKeegan et al. (2011) using the Genesis sample return mission. After correction for isotopic fractionation in the solar wind, they estimated solar isotopic ratios of 16 O / 18 O = 530 and 16 O / 17 O = 2798, which we adopt here. The most recent spectroscopic analyses of CO lines in the Sun with 3D atmosphere models are consistent with these ratios, albeit with relatively large uncertainties (Lyons et al. 2018). Like almost all rocky material in the inner Solar System, the Earth is enriched in 18 Oand 17 Oin comparison to 16 O(isotopic ratios of 499 and 2632, respectively, Meija et al. 2016b), likely as a result of CO self-shielding during the proto-planetary disc phase of the Solar System's formation (Clayton 2002; McKeegan et al. 2011).  \nGenesis solar wind data also allow the isotopic abundances of the other noble gases Ne, Ar, Kr, and Xe to be determined. For Ne and Ar, the isotopic fractionations are significantly smaller  \nthan for He, while for Kr and Xe they can be neglected altogether (Heber et al. 2012). Using the empirical correlations to correct for the fractionation, the present-day solar isotopic ratios are 20 Ne / 22 Ne = 13 : 36 GLYPH<6> 0 : 09, 21 Ne / 22 Ne = 0 : 0324, and 36 Ar / 38 Ar = 5 : 37 GLYPH<6> 0 : 03 (Heber et al. 2012; Meshik et al. 2014); almost identical values are obtained with the ine GLYPH<14> cient Coulomb drag theoretical model (Bodmer & Bochsler 2000). To our knowledge, the 36 Ar / 40 Ar ratio has not been determined from the Genesis samples and hence the proto-solar 40 Ar abundance is highly uncertain (Anders & Grevesse 1989); the terrestrial abundance is completely dominated by the radioactive decay of 40 K (Meija et al. 2016b), making it the most abundant isotope. For Kr and Xe, we embrace the Genesis results from Meshik et al. (2014) and Meshik et al. (2020), respectively.  \nFor the non-volatile elements, the Genesis solar wind samples have not yet delivered final isotopic abundances, although this should be feasible for elements like Mg, Si, S, and Fe in due course (Burnett et al. 2019). Mg is particularly interesting since it enables a critical test of the solar wind isotopic fractionation models (e.g. Bodmer & Bochsler 2000; Laming et al. 2017), which are in reasonable agreement with the most recent Genesis results (Jurewicz et al. 2020). Fe isotopic measurements of the Genesis data on the other hand would test the notion of a chemically homogeneous proto-solar nebula, that is, whether the Sun and, for example, the most chemically pristine meteorites are expected to have identical composition as normally assumed (Sect. 3).  \nFor the elements not discussed above, we take the presentday relative isotopic fractions from the recommendations by the International Union of Pure and Applied Chemistry (Meija et al. 2016b with subsequent updates for Yb, Ta, Ir, and Hf 14 ) for representative terrestrial samples (the mean when a range is given). The proto-solar isotopic fractions given in Table B.1 are assumed to be the same as the present-day values, with the exception of radioactive elements and their daughter products; we adopt the recommended isotopic half-lives from IAEA Nuclear Data Services 15 and a Solar System age of 4.568 Gyr (Bouvier & Wadhwa 2010; Connelly et al. 2012).", '5. Present-day and proto-solar Solar System abundances': 'With the Sun providing > 99 : 85% of the total mass in the Solar System, we can safely adopt the solar abundances as the basis of the chemical composition of the solar system overall. The present-day photospheric elemental abundances, however, are not the same the Sun was born with due to three main effects: nuclear processing, radio-active decay, and di GLYPH<11> usion. The first e GLYPH<11> ect only impacts Li (and possibly Be), which has been largely destroyed in the convection zone, as is evident from the comparison with the most pristine meteorites (Sect. 3.2). No nuclear products from H-burning in the solar core have reached the solar surface layers, with the exception of D-burning in the deep convection zone during the pre-main sequence (Sect. 4). Secondly, elements such as Th and U are radioactively unstable, and hence their proto-solar abundances 4.57 Gyr ago were larger than measured today; as noted in Sect. 4, we account for this using the recommended half-lives. Finally, all elements heavier than H have partly settled from the surface convection zone to the radiative interior from the combined e GLYPH<11> ects of microscopic  \nFig. 9. Present-day photospheric mass fractions the most abundant elements. Left panel: Mass fractions of H, He, and heavy elements. Right panel: Relative mass fractions of the most abundant heavy elements with the highly volatile C, N, O, and Ne contributing more than 77% of all heavy elements by mass.  \n<!-- image -->  \nthermal di GLYPH<11> usion, gravitational settling, and radiative acceleration, normally together referred to as atomic di GLYPH<11> usion (Thoul et al. 1994).  \nQuantitative predictions for elemental di GLYPH<11> usion are hampered by the possibility of additional mixing below the convection zone due to still poorly understood processes like convective overshooting, rotation, internal gravity waves, and turbulence (e.g. Goldreich & Schubert 1967; Chaboyer et al. 1995; Bara GLYPH<11> e et al. 2017; Aerts et al. 2019; Dumont et al. 2021), which is required to explain the solar Li abundance and interior rotation profile. Di GLYPH<11> erent groups thus require slightly varying proto-solar abundances depending on the ingredients of their stellar evolution codes when calibrating to the solar luminosity, e GLYPH<11> ective temperature and surface composition at the solar age. Here we adopt the predictions from the standard solar model of Vinyoles et al. (2017) computed with the Asplund et al. (2009) chemical composition. In terms of number densities relative to H, the proto-solar He abundance is higher by 0.070 dex compared with the present-day photospheric value, while for all heavier elements the di GLYPH<11> erence is 0.064 dex. Because the Vinyoles et al. (2017) solar model does not include any additional mixing beyond atomic di GLYPH<11> usion as required to explain the solar Li depletion, these di GLYPH<11> usion corrections may be slightly over-estimated; more recent solar models which simultaneously agree better with the Li depletion and helioseismology (see below) have di GLYPH<11> usion abundance corrections about half of this: GLYPH<25> 0 : 03 dex for all elements other than H (Eggenberger et al. 2021).  \nRegarding elements for which no photospheric abundance determination is feasible from spectroscopy or solar wind measurements (Table 2), we base their present-day photospheric abundances on the CI chondritic abundances. Those are first corrected for the trend in abundance di GLYPH<11> erence with condensation temperature (Fig. 6, Sect. 3.3) before applying the e GLYPH<11> ects of atomic di GLYPH<11> usion to compute the proto-solar abundances. The proto-solar isotopic abundances are given in Table B.1.  \nIn terms of mass fractions of H, He, and heavy elements, the elemental abundances presented in Table 2 correspond to present-day photospheric values of X surface = 0 : 7438 GLYPH<6> 0 : 0054,  \nY surface = 0 : 2423 GLYPH<6> 0 : 0054, Z surface = 0 : 0139 GLYPH<6> 0 : 0006, and Z surface = X surface = 0 : 0187 GLYPH<6> 0 : 0009; Fig. 9 illustrates the presentday photospheric mass fractions of the most abundant elements. The quoted uncertainty on Y surface stems from the helioseismic determination (Sect. 2.8), while for the heavier elements we assumed that their abundance errors to be independent. The here inferred solar surface metallicity is slightly higher than (but within the uncertainties of) the value ( Z surface = 0 : 0134) derived in Asplund et al. (2009), mainly as a result of the revised Ne abundance following the adoption of the Genesis solar wind measurements (Sect. 2.8). With the Vinyoles et al. (2017) solar model for atomic di GLYPH<11> usion, the proto-solar mass fractions would be X initial = 0 : 7121, Y initial = 0 : 2725, and Z initial = 0 : 0154 ; while with a smaller abundance correction of 0.03 dex for all elements heavier than H (Eggenberger et al. 2021) the values are X initial = 0 : 7304, Y initial = 0 : 2550, and Z initial = 0 : 0146. Table 4 lists the inferred present-day photospheric and proto-solar mass fractions based on some widely used compilations of the solar chemical composition over the past three decades.', '6. Solar abundances, standard solar models, and helioseismology': 'Care must be exercised when interpreting proto-solar mass fractions in light of the long-standing and exasperating discrepancy between helioseismology and predictions from standard solar models constructed with present-day surface abundances similar to those presented here: the so-called solar modelling problem 16 (e.g. Basu & Antia 2004; Bahcall et al. 2005; Vinyoles et al. 2017; Christensen-Dalsgaard 2021). In particular, the comparatively low photospheric C, N, and O abundances (and associated Ne abundance) from 3D-based spectroscopic analyses (e.g. Allende Prieto et al. 2001b, 2002; Asplund et al. 2004, 2005b, 2009; Amarsi et al. 2018a, 2019, 2020) have ruined the good agreements with helioseismology that were obtained with  \nTable 4. Present-day photospheric and proto-solar mass fractions of hydrogen (X), helium (Y), and metals (Z) for a number of widely used compilations of the solar chemical composition.  \nNotes. ( a ) Ca GLYPH<11> au et al. (2011) determined the abundances of Li, C, N, O, P, S, K, Fe, Eu, Hf, Os, and Th and adopted the recommended values of Lodders et al. (2009) for all other elements. They did not quote any proto-solar mass fractions.  \nglobal solar models computed with the older higher solar metallicity (e.g. Anders & Grevesse 1989: Z surface = X surface = 0 : 274; Grevesse & Noels 1993: Z surface = X surface = 0 : 246; Grevesse & Sauval 1998: Z surface = X surface = 0 : 231). The predicted surface He abundance, depth of the convection zone, and the sound speed as a function of solar radius from a solar calibration are now significantly more discrepant compared to the inversion of the observed solar p-mode oscillation frequencies; the solar surface composition presented here goes in the right direction but will not resolve the mismatch. We note, however, that the available seismic data in fact favours a low metallicity ( Z surface < 0 : 014) in the convection zone (Vorontsov et al. 2013; Buldgen et al. 2017), suggesting that current standard solar models are missing important physics or are constructed with faulty input data in the radiative zone below the convection zone.  \nA huge body of work has been devoted to reconciling the spectroscopic and interior solar models triggered by our 3Dbased solar abundances over the past two decades. Arguably the most straightforward explanation would be underestimated opacities, especially near the bottom of the convection zone where the discrepancies are most acute, to compensate for the lowering of the abundances of the dominant elements. It would require the computed data sets of atomic opacities (e.g. Badnell et al. 2005; Mondet et al. 2015; Colgan et al. 2016) all to be underestimated by about 20% in the relevant temperature-pressure regime (Serenelli et al. 2009; Christensen-Dalsgaard et al. 2009). In fact recent experimental opacity measurements are systematically higher than calculations with the missing opacity for Fe alone accounting for about half of what is required to resolve the solar modelling problem (Bailey et al. 2015; Nagayama et al. 2019). With new experiments and updated calculations with previously overlooked atomic processes underway (e.g. Nahar & Pradhan 2016; Krief et al. 2018) it remains to be seen whether the final atomic opacities will be su GLYPH<14> cient. Important independent constraints on the solar interior composition will also come from measurements of the solar neutrino fluxes, especially those produced in the CNO-cycle that were recently reported by the Borexino experiment (Borexino Collaboration et al. 2020). Finally, encouraging signs are appearing from new solar models with improved mixing below the convection zone aimed at explaining the depletion of Li and the interior rotation profile, which standard models cannot do (e.g. Yang 2019; Zhang et al. 2019; Buldgen et al. 2019b,a; Eggenberger et al. 2021). We re-  \nmain cautiously optimistic that the persistent solar modelling problem will soon have been resolved without necessitating any major revision of the photospheric solar abundances presented here.', '7. Concluding remarks': 'The elemental abundances derived and presented here are a further stepping stone towards the advancement of our understanding of the solar chemical composition, which is of such crucial importance for astronomy as a whole. As demonstrated above, major improvements in terms of the solar spectroscopic analysis have been accomplished, which should make the inferred elemental abundances both more precise and more accurate. It is now perfectly feasible to perform detailed 3D non-LTE line formation calculations using highly realistic 3D radiativehydrodynamical simulations of the solar surface convection and atmosphere. This has been achieved for 13 elements - Li, C, N, O, Na, Mg, Al, Si, K, Ca, Mn, Fe, and Ba - all using comprehensive model atoms with reliable atomic data, including the previously notoriously uncertain cross-sections for H collisions. Compared to Asplund et al. (2009), in which only the abundances of Li and O were based on a consistent 3D non-LTE analysis, this certainly represents substantial progress. The abundances of a further 45 elements have been determined on the basis of a 3D LTE spectroscopic analysis. In total, we present the solar photospheric abundances of 68 elements when also including the noble gases, for which helioseismic and solar wind measurements are available, and four trace elements evaluated from sunspot observations. Importantly, the abundance uncertainties for these 68 elements have been carefully evaluated, including both statistical and systematic uncertainties. These have been augmented with meteoritic data (corrected for an identified trend with condensation temperature) for a further 15 elements, thus providing Solar System abundances for all 83 long-lived elements.  \nBased on our state-of-the-art 3D non-LTE analysis, we can safely confirm that the old high solar metallicity of, for example, Grevesse & Noels (1993) and Grevesse & Sauval (1998), which are still used to construct standard solar models, should finally be discarded - there are simply no arguments in favour of those outdated 1D LTE results spectroscopically. Our study is, however, definitely not the final word on the issue of the so-  \nlar chemical composition, with much work remaining to reach comparable abundance precision to that of the most pristine meteorites. Arguably, the most pressing need is to extend non-LTE calculations to additional elements: only 26 elements have their solar abundances determined from a non-LTE analysis, half of which are only in 1D. Of particularly high priority is to carry out 3D non-LTE computations for the remaining Fe-peak elements (including revisiting the important case of Mn, Sect. 2.5) as well as other elements of key astronomical significance, such as Be (to constrain solar depletion, Sect. 3.2), S, Zn, Rb (volatile elements to confirm di GLYPH<11> erences with CI chondrites, Sect. 3.3), P, Cu, and Sr (stellar nucleosynthesis yields, Kobayashi et al. 2020). Several of these elements already show substantial nonLTE e GLYPH<11> ects in existing 1D non-LTE calculations, which can be expected to be amplified in 3D.  \nAs always, solar spectroscopy is reliant on having access to accurate atomic data. Over the past decade, the situation regarding experimental transition probabilities has improved significantly for the Fe-peak elements in particular, thanks to the dedicated and valiant e GLYPH<11> orts by a relatively small group of atomic physicists around the world, as discussed in Sect. 2.5. Several other elements have also seen significant updates, however, the need for further work is essentially endless. Here it su GLYPH<14> ces to mention the case of Si, the element anchoring the meteoritic abundance scale to the solar photospheric values (Sect. 2.4). Besides the need for refined transition probabilities of spectral lines used as diagnostics, there is also a demand for more and better atomic data for lines blending the primary abundance indicators. Since those lines are typically weaker and arise from higherexcitation levels, they are often not even properly identified or, when included in existing atomic calculations, have faulty gf -values.  \nFor non-LTE modelling, significantly more atomic data are necessary, including all radiative (bound-bound and bound-free) and collisional (with electrons and hydrogen) transitions coupling the di GLYPH<11> erent atomic levels. Great progress has recently been achieved in this regard, especially related to inelastic H collisions (e.g. Barklem 2016a; Belyaev et al. 2018; Barklem et al. 2021), which has made the non-LTE results much more trustworthy. Crucially, the relevant atomic data now exist for a 3D non-LTE study for almost all of the elements highlighted above; mainly what is lacking is (wo)man-power to take on the necessary but time-consuming 3D non-LTE analyses. What has not yet been fully explored for the Sun is how departures from LTE in one element a GLYPH<11> ect those of other elements or feed back onto the solar atmospheric structure; preliminary calculations suggest that this may well be of importance (Osorio et al. 2020). Furthermore, we caution that magnetic fields have not been included in the 3D solar atmosphere model employed throughout this study. While early work based on imposed vertical magnetic fields implied that the e GLYPH<11> ects could be substantial (Fabbian et al. 2010), more recent 3D LTE calculations based on a 3D MHD solar model with a self-consistent small-scale magnetic dynamo (Rempel 2014) reveal that the impact on solar abundance determination is negligible ( < 0 : 01 dex) at least for C, N, O, and Fe (Shchukina & Trujillo Bueno 2015; Shchukina et al. 2016); this may not hold however for elements in which non-LTE e GLYPH<11> ects are important and thus deserves further scrutiny. Similarly, there is no indication that the inclusion of a chromosphere or other magnetic activity in the solar model would significantly impact the derived abundances for the weak photospheric lines employed here using disc-centre observations of the quiet Sun (Bergemann et al. 2021).  \nBesides the obvious benefits from having a better definition of a widely used astronomical yardstick, the continuously improved 3D-based solar abundances reported here and in our earlier works (e.g. Asplund et al. 2005a, 2009) have uncovered two surprising peculiarities regarding the solar chemical composition, neither of which have been explained. The inconsistency between the composition at the surface, as measured from spectroscopy, and the interior, as inferred from helioseismology, has now been around for almost two decades in spite of tremendous e GLYPH<11> orts in reconciling the two with further improvements to the solar modelling. It is crucial to remember that this is not a solar problem but a stellar problem . If indeed the explanation is missing opacity as suggested by recent experiments (Sect. 6), it would imply that current stellar interior and evolution models are not fully trustworthy. Since such stellar models underpin much of the field of astronomy, it would have profound implications. The second solar abundance feature uncovered here, namely that the supposedly most primitive meteorites - the CI chondrites and the Sun in fact do not share the same composition, implies that a revision of the earliest history of the Solar System is required, given that their equivalence has been taken for granted for more than half a century (Sect. 3.3). The existence of a similar abundance trend of the condensation temperature between the Sun and solar twins suggests that this may well give us a deeper understanding of planet formation around stars in general.  \nAcknowledgements. We thank all of our wonderful colleagues around the world for highly rewarding collaborations over many years, numerous stimulating discussions about the topics described herein and their implications for astronomy and physics, and for providing necessary input atomic / molecular data, all of which have greatly benefited this work. MA gratefully acknowledges generous funding from the Australian Research Council through a Laureate Fellowship (FL110100012) and a Discovery Project (DP150100250). AMA acknowledges support from the Swedish Research Council (VR 2016-03765 and 2020-03940). This research was supported by computational resources provided by the Australian Government through the National Computational Infrastructure (NCI) under the National Computational Merit Allocation Scheme and the ANU Merit Allocation Scheme (project y89). 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Line list for the Fe lines used for 3D non-LTE spectral line formation calculations to derive the solar photospheric abundances', 'Appendix B: Isotopic abundances': 'Table B.1. Representative isotopic abundance fractions and logarithmic abundances by number (log GLYPH<15> H GLYPH<17> 12 : 00) at the time of birth of the Sun. The proto-solar abundances account for atomic di GLYPH<11> usion of 0 : 07 dex for He and 0 : 064 dex for all heavier elements as predicted by Vinyoles et al. (2017), as well as radio-active decay. The elements for which no solar abundance is available are based on CI chondrites adjusted for the trend in condensation temperature between the Sun and meteorities (Sect. 3.3).'}